Planetary Science Mars Research 1

PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

posted July 1, 2004

Carbonates in ALH 84001: Part of the Story of Water on Mars --- The study of multi-generational carbonate assemblages in Martian meteorite ALH 84001 reveals a complex history of crystal formation, growth, and alteration.

Written by Catherine M. Corrigan Smithsonian Institution, National Museum of Natural History

Carbonate-rich regions in ALH 84001 are complicated. There are familiar forms of carbonate as well as fascinating textural forms previously unreported including carbonate rosettes, planiform "slab" carbonates, distinct "post-slab" magnesium carbonates (magnesite), and carbonates interstitial to feldspathic glass and orthopyroxene. Slab carbonates reveal portions of the carbonate growth sequence not seen in the rosettes and suggest that initial nucleating compositions were rich in calcium. They formed in two major stages. The first stage involved growth of the rosettes and slab carbonates. This step was controlled by the rate of crystal nucleation, how fast the ingredients were delivered to the growing crystals, and how much fluid was available. Cosmochemists call this type of growth "kinetically controlled." Next, an alteration event formed the magnesite-siderite (iron carbonate) layers on the exterior surfaces of the carbonate. Post-slab magnesite, intimately associated with silica glass, is compositionally similar to the magnesite in these secondary exterior layers, but represents a later generation of carbonate growth. Formation of feldspathic glasses had little or no thermal effect on carbonates, as indicated by the lack of thermal decomposition or any compositional changes associated with glass/carbonate contacts. The carbonates tell an important story about water in the ancient crust of Mars. The presence of numerous, distinct generations of carbonate formation and relatively clear fracture chronology within carbonate further suggest that interactions between ALH 84001 and the crustal fluids of Mars were discontinuous and occurred only a few times over its 4.5 Ga history. The reactivation and remobilization of fluids (causing events such as formation of magnesite-siderite-magnesite layers and precipitation of post-slab magnesite) and the fracturing within the rock were almost certainly driven by impacts. The evidence for punctuated, impact-driven interaction between rocks and fluids supports scenarios describing temporary hydrous environments as opposed to those including large-scale, long-term hydrologic systems including oceans. Therefore, unless ALH 84001 is a particularly rare sample that escaped intense alteration, the hydrosphere of Mars may not have interacted with the rocks as thoroughly as planetary geologists have inferred from the presence of river networks and other features formed by flowing water. 1 of 15

PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

Reference: Corrigan C. M. and Harvey R. P. (2004) Multi-generational carbonate assemblages in Martian meteorite Allan Hills 84001: Implications for nucleation, growth and alteration. Meteoritics and Planetary Science, v. 39, p. 17-30.

Water on Mars

A central goal of the Mars Exploration Program is to understand the history of water on Mars. We need to know how much there is now and was soon after the planet formed, how it cycles through the crust, where it resides now, how much is in magmas produced in the mantle and how it varies throughout the mantle, and how it has shaped the surface of the planet. The quest to understand all about water on Mars is part of a larger effort to determine if the planet was ever habitable enough for life to have originated and evolved. (Detailed goals for the exploration of Mars have been developed by the Mars Exploration Program Analysis Group, MEPAG, and can be found at the MEPAG web site. Link opens in a new window.) Planetary geologists have identified numerous features that indicate that vast amounts of water sculpted the Martian surface: valley networks, huge outflow channels, layered sediments, and recent gullies. There might even have been an ocean in the northern lowlands of Mars. Recent observations by the Opportunity rover in Meridiani Planum show that water reworked sediments and deposited a sequence of minerals as it evaporated. Cosmochemists have calculated that the salty water that made all these impressive features would have evaporated to produce vast quantities of carbonate minerals. However, observations by the Thermal Emission Spectrometer onboard the Mars Global Surveyor spacecraft indicate that there is only a few percent of carbonate on the surface. Martian meteorites contain a little carbonate, which studies show formed on Mars, not after their arrival on Earth, but the amount is very small. Clearly we are missing something important. Martian meteorite ALH 84001 has the most carbonate of any Martian meteorite, so it might hold the key to understanding carbonate formation on Mars. The carbonates in ALH 84001 are an important part of the story of water on Mars.

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

LEFT: These layers of rock were probably deposited by water on Mars. They look very similar to layers of sedimentary rock in many parts of the world, such as the western United States. [More information on this MOC image from Malin Space Science Systems.] RIGHT: The dendritic pattern of these fine channels and their location on steep slopes suggest they are runoff channels. [More information on this Viking image from NASA/JPL Planetary Photojournal.]

An Important but Confusing Meteorite

Following its identification in 1994 as a Martian meteorite by David Mittlefehldt (Johnson Space Center, JSC), Allan Hills (ALH) 84001 was rapidly recognized as being a very important sample of the Martian crust (pictured on the right). Its fame exploded with the publication of a paper in 1996 by David McKay and his colleagues from JSC, Stanford University, and McGill University that said the meteorite contained evidence for life on Mars. [See PSRD article Life on Mars?] This announcement led to a huge number of studies designed to test this interpretation. [See the PSRD archive of articles on Mars Life Issues.] Whether it contains fossil life or not, ALH 84001 was affected by water-bearing fluids while it was still home in the ancient crust of Mars. ALH 84001 is by far the oldest of the Martian meteorites (~4.5 Ga). ALH 84001 is a cumulate orthopyroxenite, which means it is loaded with a mineral called orthopyroxene. The orthopyroxene is accompanied by some chromite, feldspathic glass, augite, apatite and olivine. It formed when a magma invaded the juvenile Martian crust. As the magma crystallized, orthopyroxene accumulated in the magma. Slow cooling in a big magma body allowed time to make big crystals, so it is coarse grained. However, it has a cataclastic texture, indicating that it has been exposed to a series of impact shock events that partly demolished the original igneous texture. This intense shock metamorphism has resulted in the presence of crushed zones, or granular bands, that contain crushed orthopyroxene, chromite, feldspathic glass, olivine and other phases, including the all-important carbonates that Ralph Harvey (Case Western Reserve University) and I studied. Despite being one of the most studied rocks of all time, we do not understand its complex 3 of 15

PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

history. Many distinct geologic stories have been told for ALH 84001. Most agree on initial crystallization as part of a slow-cooling underground magma chamber, and that the rock gained its current, highly fractured state during several post-crystallization impact events. The time between impacts is not known. The resultant fractures provided conduits for the passage of fluids through the rock, and allowed the development of secondary, non-igneous minerals within them. The secondary minerals are of particular interest in ALH 84001 because they offer physical and chemical clues to past Martian environments. Carbonate minerals are the dominant secondary phase in ALH 84001, making up ~1% of the rock. They occur in a variety of settings and textures, from interstitial crack fillings to conspicuously zoned clusters, semi-circular in cross-section, which have gotten the name "rosettes." ALH 84001 rosettes vary in composition concentrically from Ca-rich near the center through dolomite-ankerite (Mg-rich) to alternating magnesite-siderite-magnesite (MSM) layers at the outer edges, the siderite layers of which often contain fine, single-domain magnetite, a central part of the debate about evidence for life in the meteorite. The ages of the carbonates are difficult to determine, but careful work by Grenville Turner (University of Manchester) and colleagues and Lars Borg (University of New Mexico) and colleages indicate an age of 3.83 Ga to ~4.04 Ga, coincident with the period of heavy bombardment in the inner solar system prior to ~3.8 Ga. On the basis of their textural setting in the rock and on their ages, it is safe to conclude that carbonate formation clearly post-dates both initial igneous crystallization and an initial episode of fracturing by impact. Cosmochemists are still debating how the carbonates formed. Proposed formation scenarios include low-temperature aqueous precipitation, evaporative processes, high-temperature reactions, and impact-induced melting. Recent experimental studies by D. C. Golden (Hernandez Engineering Inc., Houston) and colleagues at JSC confirm that low-temperature precipitation (150 oC) from a saturated fluid followed by short-term heating can reproduce many of the carbonate features seen in ALH 84001. Uncertainty also lies in whether there were single or multiple generations of carbonate formation and the role, if any, of alteration after it had formed. Adrian Brearley (University of New Mexico) suggested that some nanometer-scale mineralogy and textures in ALH 84001 result from thermal decomposition of pre-existing carbonate materials, particularly siderite decomposition to magnetite. The experiments of D. C. Golden and colleagues showed that subsequent heating of carbonates formed in the laboratory (to 470 oC) was adequate to form magnetite crystals. Alternatively, Kathie Thomas-Keprta (Lockheed Martin) and her colleagues at JSC still believe that at least a subset of the tiny crystals of magnetite in ALH 84001 were made by Martian microorganisms. The carbonate minerals in ALH 84001 record part of the story about water in the ancient Martian crust, but the complexity of the carbonates and the drastically different interpretations of how they formed obscure the story. Ralph Harvey and I tried to clear up the confusion by studying several complex, carbonate-rich regions in ALH 84001. We examined forms of

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

carbonate familiar to cosmochemists who had studied ALH 84001, as well as more complete exposures of carbonate growth sequences. We placed our observations into context with previous work, offering insight into the complicated story of carbonate formation in this unique and important meteorite.

What the Carbonates Look Like

The regions examined in this study revealed a variety of textural relationships ranging from simple to complex, with carbonate, feldspathic, and silica glasses being the most significant phases. Carbonate occurs as the commonly described, photogenic rosettes, and in three other forms seen previously but not described or classified in a uniform way. These are discrete, layered packages here termed "slab" carbonates, massive background fill ("post-slab magnesites") and carbonate occurring "interstitial" to feldspathic glasses and orthopyroxene.

Backscattered electron image of a region of secondary minerals in ALH 84001,303 (the ",303" indicates the subsample number). This image shows carbonate rosettes (CR) to the left, feldspathic glass (Fs), slab carbonates (Slab), interstitial carbonate (IC), and orthopyroxene (Opx).

Rosettes. Rosettes were found only in ALH 84001,303 (see image above). These rosettes are identical to those described and analyzed in many previous studies of ALH 84001, with semi-circular cross-sections and distinct, consistent and concentric chemical zoning. This

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

zoning includes alternating magnesite-siderite-magnesite (MSM) layers on the outer portions of the carbonate sequence.

Backscattered electron image of a region of secondary minerals in ALH 84001,302. Slab carbonates (Slab) are shown. Dark, Mg-rich carbonates found between feldspathic glass and slabs are termed post-slab magnesite (PSM) and are often mixed with silica glass (Si). Feldspathic glass (Fs), interstitial carbonate (IC), and orthopyroxene (Opx) are also shown.

Slab. We identified a previously undescribed form of carbonate, which we termed "slab" carbonate. It figures prominently in both of the regions studied in detail. These slab carbonates are elongate packages that conform to fracture faces and exhibit chemical zoning distinctly visible in backscattered electron images taken with an electron microprobe, duplicating zoning commonly seen in rosettes. (See images above.) This zoning, however, is parallel to the faces of the slab (instead of concentric around a central point) and exhibits a more complete chemical zoning record with sharper contacts between compositional zones. Slab carbonates include a thin, Ca-rich layer (bright in the electron images) at one edge and the familiar magnesite-siderite-magnesite (MSM) layers at the opposite edge, with three consistently distinct layers found between. Slab carbonates are typically highly fractured, in a manner indistinguishable from that seen in rosettes, with fractures generally crossing all layers. (See image below.) Like rosettes, slabs show no obvious preferential association with specific mineral species. They are found in contact with orthopyroxene, feldspathic glass, and other carbonates.

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

Backscattered electron image of ALH 84001,303 showing slab carbonate (Slab), post-slab magnesites (PSM), silica glass (Si) and magnesite-siderite-magnesite (MSM) rims. Orthopyroxene (Opx) is present at the top of the image. Interstitial carbonates (IC) can be seen mixed with both orthopyroxene and feldspathic glass (Fs).

Post-slab magnesite. The regions studied here contain almost pure magnesite that is texturally distinct from carbonate in rosettes or slabs. We refer to these carbonates as "post-slab magnesites" as we believe they formed as a distinct generation, post-dating zoned slabs, rosettes, and magnesite-siderite-magnesite (MSM) layers. Post-slab magnesites occur as numerous blebs or grains, semi-circular in cross section (see "PSM" in the images above). They do not appear to be the outer edges of rosettes, and are smaller and more uniform in size. These carbonates entrain small fragments of other minerals, including other carbonates. Post-slab magnesites have a fracturing habit different from the zoned carbonates, with fractures formed around individual blebs, and rarely crossing through them. In the regions we studied, we found post-slab magnesite in contact with the high-Ca end of the slab carbonate (bright in the photographs). It occurs with silica glass and mixed with fragments of other phases. Silica glass tends to be associated with fractures surrounding post-slab magnesite blebs, and occasionally occurs as larger irregular blebs (see image below).

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

Backscattered electron image of post-slab magnesite (PSM) in ALH 84001,302 showing blebby texture of carbonate, silica glass (Si) rims around black carbonate blebs, and larger blebs of silica glass. This image was contrast enhanced to better show the mottled texture within the post-slab magnesite region. Slab carbonate (Slab), orthopyroxene (Opx), and feldspathic glass (Fs) are also labeled.

Interstitial carbonate. We treated carbonate interstitial to larger domains of other minerals, particularly feldspathic glass and comminuted orthopyroxene, as a separate category called interstitial carbonate (see "IC" in the images above). Some carbonates appear to have been mechanically entrained by feldspathic glass, which accounts for a significant portion of the mineral inventory in the regions we studied. Carbonates found in the interstices of orthopyroxene surrounding the secondary mineral regions are similar in appearance to those found entrained within feldspathic glasses. The carbonates in orthopyroxene have been described previously.

What the Carbonates are Made Of

I measured the composition of the carbonates using an electron microprobe. The diagram below shows ~800 carbonate analyses from these sections. The data are plotted on a triangular diagram, called a ternary diagram. Each corner of the triangle represents the pure composition of a mineral. All data are plotted as molar abundances of MgCO3, CaCO3 and FeCO3. The distance from a corner gives the abundance of that chemical component in a mixed mineral.

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

Ternary diagram of major element compositions of all carbonates I measured in ALH 84001,302 and ALH 84001,303 and individual textural groups.

My new analyses overlap and extend compositional ranges previously reported for ALH 84001 carbonates. These new data show a much more continuous compositional trend filling gaps seen in previous work, including significant proportions of high-Ca carbonate that were seen only sporadically in previous studies.

Ternary diagrams of major element compositions of carbonate rosettes (left) and slabs (right). Compositions from the interiors of rosettes (more central to the ternary) are clearly distinguishable from those in the magnesite-siderite-magnesite (MSM) rims (seen at the MgCO 3 apex). Compositions of slab carbonates span the entire range of compositions.

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

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Compositions of slab carbonates differ from rosettes only in that they exhibit a wider range of values, varying in a nearly continuous sequence across the ternary compositional diagram (see diagrams above). The sharp boundaries between layers seen in backscattered electron images also correspond to measured compositional changes. Point analysis transects were constructed approximately perpendicular to zoning across slab carbonates in one of the samples studied and ALH 84001,302 revealing a consistent sequence of compositional variation. This sequence suggests that the compositional sequence seen in carbonate rosettes is a subset of that seen in slab carbonates.

Ternary diagrams of major element compositions of post-slab magnesites (left) and interstitial carbonates (right). As their name implies, post-slab magnesites cluster in the Mg-carbonate apex of the diagram, though they do reach farther toward intermediate compositions than do the MSM rims seen in rosettes. Compositions of interstitial carbonates span the entire range of compositions.

The massive, space-filling post-slab magnesites span a wider compositional range than magnesites in the magnesite-siderite-magnesite layers of rosettes, from nearly pure MgCO3 to intermediate compositions. The interstitial carbonates are not chemically distinct. Their compositions span nearly the entire range, though most are intermediate, suggesting that they represent a combination of all observed carbonate sources (see diagrams above).

How the Carbonates Formed

The concentric zoning of the rosettes suggests that they formed in small pockets of water by nucleating (forming a seed crystal) in one place and growing from it to fill up the space until the water was used up. As the rosettes grew, their formation changed the composition of the surrounding water, causing them to become chemically zoned. They nucleated on any other crystal (large orthopyroxenes, crushed orthopyroxene, feldspar) and do not appear to have reacted with the other minerals. This type of crystal growth producing zoned crystals is called 10 of 15

PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

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kinetically controlled--the growth is governed by the abundance and speed of delivery of the raw materials to the growing carbonate crystals. The textures of microcrystalline slab carbonates are also consistent with formation under kinetically controlled conditions. Slab carbonates show visible zoning parallel to fracture surfaces, suggesting nucleation from numerous, closely spaced points on a surface instead of from widely spaced individual points. The identical zoning in rosettes and slabs suggests that they formed during the same growth event; the difference between the two forms relates to the amount of available space. Essentially, when there is sufficient fluid volume (i.e, in larger fractures) slabs will form, while rosettes will form when volume is limited. Slab carbonates are of particular value in understanding the sequence of carbonate crystallization. Their semi-planar geometry offers an advantage in that a random slice through a slab is more likely to intersect the full range of compositions present in three dimensions than is a slice through a rosette. Slab carbonates should thus provide a more complete history of carbonate formation, exposing compositions representative of early stages of formation (high-Ca layers) rarely seen in rosettes.

A Five-Step Sequence of Events

Our observations lead us to propose the history of carbonate formation shown in the diagram below. Initial formation of the rock as a cumulate orthopyroxenite was followed by impact events resulting in an initial set of fractures within the rock. A carbonate growth stage occurred next, during which rosettes and slab carbonates were precipitated into the fractures from water super-saturated in carbonate components. We do not know where the water came from, but it seems clear that the rock and its surroundings were not saturated in water for a long time because of the low abundance of carbonates and very limited alteration of the original igneous minerals. Rosettes arose from isolated nucleation sites in relatively small fractures, forming pancakes where perpendicular growth was limited and more spheroidal shapes where space allowed. Slabs nucleated in rare, larger fractures. Earliest formed carbonates were Ca-rich, but crystallization progressed toward more Mg- and Fe-rich compositions. Occasional recharge of fluids during carbonate growth altered the cation concentrations resulting in the variable compositions visible as zones in the backscattered electron images shown above.

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html

The chemistry, mineralogy, and textures of carbonate lead us to favor a thermal event as the mechanism for forming the magnesite-siderite-magnesite (MSM) layers (as suggested by Adrian Brearley of the University of New Mexico, and others) rather than a dramatic change in the chemical composition of the fluid. We suggest that this event caused the conversion of an exterior, Mg- and Fe-rich carbonate composition into the residual MSM layers, which may have involved replacement or dissolution and redeposition of carbonate materials. The original carbonate growth trend from high to lower Ca compositions (slabs) was replaced with carbonates of relatively constant Ca and widely varying proportions of Fe and Mg (MSM). This 12 of 15

PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

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event would have occurred with the earliest slab carbonates still attached to nucleation surfaces, as these innermost compositions appear to be unaffected by alteration. In addition, MSM layers are found concentric only around the exteriors of zoned carbonates, which would have required free space between these surfaces and the fracture walls. Textures suggest that post-slab magnesites represent a distinct precipitation event that took place after the MSM sequences formed and before feldspathic glass intruded. Most fractures crossing slab carbonates and MSM layers do not cross into post-slab magnesites. Entrainment of slab carbonate material into post-slab magnesites also suggests that these magnesites formed after rosettes and slabs. Post-slab magnesites are physically mixed with silica glasses, suggesting that either the two phases were emplaced during the same event or silica glass precipitated after post-slab magnesite. Unlike MSM bands, post-slab magnesites are found in contact with the oldest (earliest formed) slab surfaces. The Ca-rich edges of the slab must have been detached from their original nucleation surfaces allowing space for post-slab magnesite to precipitate. Chemically, MSM magnesites and the post-slab magnesite are similar suggesting that they formed either by similar processes or that they represent two stages of a single event. The first decomposed existing carbonates and deposited the MSM rims, while the second precipitated the chemically similar post-slab magnesites and silica glass, filling in the trend with compositions lost during the first event. The final step involves the injection of feldspathic glass. There are numerous occurrences of zoned carbonates and post-slab magnesite entrained by feldspathic glass. In addition, there exist locations where fractures transcend the boundaries between both types of carbonate but do not cross into feldspathic glasses. These observations indicate that feldspathic glass was the last phase to enter these fractures. The bulbous texture and lobate contacts with post-slab magnesite suggest that feldspathic glass was mobile and flowed into these fractures. Feldspathic glass intrusion produced physical effects but did not seem to cause significant chemical changes. Glass intrusion further widened the fractures, entraining phases already present in fractures and further peeling some carbonates from their nucleation sites. The intrusion of feldspathic glass separated the slabs in ALH 84001,303, as well as the attached post-slab magnesites. Other investigators have suggested that injection of the feldspathic glass caused heating and decarbonation of pre-existing carbonate. Our observations contradict this interpretation. The feldspathic glass could not have caused the heating because the sequence of events we deduced from our study indicates that the glass formed last. It was not around when the MSM layers formed. MSM sequences are present around the exterior of zoned carbonates, but neither slabs nor rosettes are consistently altered everywhere they are in contact with feldspathic glass, contrary to what would be expected if feldspathic glass was responsible for MSM formation. The variation in composition of carbonates interstitial to feldspathic glass is strong evidence that they are unaltered, mechanically entrained materials, as opposed to post-intrusion precipitates. In addition, many occurrences of slab carbonates are not visibly in contact with feldspathic glass, yet still have MSM sequences.

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

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Why did the injection of hot feldspathic glass not cause extensive heating? Studies of glass rheology may provide a solution to this paradox--it might not have been hot. Although most researchers studying ALH 84001 assume that mobilization of feldspathic glass requires high temperatures, experiments have shown that when silica glasses are exposed to high static pressures (more than 10 times the pressure at the surface of the Earth) their viscosity (resistance to flow) can drop many orders of magnitude without significant temperature elevation . Subsequent shear then easily deforms the glass with little thermal consequence. Impact events generating 45-60 GPa of pressure provide more than enough stress needed to reach this transition. As a result, feldspathic glass can flow on the millimeter scale (as suggested by textures seen in ALH 84001) in the absence of a thermal pulse. The impact event(s) that mobilized the feldspathic glasses seen in the regions studied here likely provided enough shear strength to allow the glass to flow into the fractures at low temperatures.

Implications for the Ancient Martian Crust

The presence of numerous, distinct generations of carbonate formation and relatively clear fracture chronology within carbonate show that interactions between ALH 84001 and crustal fluids on Mars were discontinuous and occurred only a few times over its 4.5 Ga history. The reactivation and remobilization of fluids (causing events such as MSM formation and precipitation of post-slab magnesite) and the fracturing within the rock were almost certainly caused by impact. The evidence for punctuated, impact-driven interaction between rocks and fluids supports scenarios describing temporary hydrous environments as opposed to those including large-scale, long-term hydrologic systems including oceans. Therefore, unless ALH 84001 is a particularly rare, particularly pristine sample, the hydrosphere of Mars may not have interacted with the rocks as thoroughly as planetary geologists infer for Mars. Geologists see clear evidence for not only river networks, but for erosion of them. Such a warm, wet period could have pervasively altered rocks in the ancient highlands, yet ALH 84001 was clearly not significantly affected by a period in which Mars was warm and wet. This could mean that ALH 84001 is just a lucky survivor. The inconsistency between photogeological and rock data needs to be reconciled before we understand the details of the history of water on Mars.

Borg, L. E., Connelly, J. N., Nyquist, L. E., Shih, C. Y., Wiesmann, H. and Reese, Y. (1999) The age of the carbonates in Martian meteorite ALH 84001. Science, v. 286, p. 90-94. Brearley, A. J. (1998) Magnetite in ALH 84001: Product of the decomposition of ferroan carbonate. Lunar and Planetary Science Conference XXIX (abstract 1451).

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PSRD:: Carbonates in ALH 84001: Part of the Story of Water on Mars

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Corrigan, C. M. and Harvey, R. P. (2004) Multi-generational carbonate assemblages in Martian meteorite Allan Hills 84001: Implications for nucleation, growth and alteration. Meteoritics and Planetary Science, v. 39, p. 17-30. Golden, D. C., Ming, D. W., Schwandt, C. S., Morris, R. V., Yang, V., and Lofgren, G. E. (2000) An experimental study on kinetically-driven precipitation of calcium-magnesium-iron carbonates from solution: Implications for low-temperature formation of carbonates in Martian meteorite Allan Hills 84001. Meteoritics and Planetary Science, v. 35, p.457-465. Golden, D. C., Lauer, H. V., Jr., Lofgren, G. E., McKay, G. A., Ming, D. W., Morris, R. V., Schwandt, C. S. and Socki, R. A. (2001) A simple inorganic process for formation of carbonates, magnetite, and sulfides in Martian meteorite ALH84001. American Mineralogist, v. 86, p. 370-375. Mazurin, O. V., Startsev, Y. K., and Stoljar, S. V. (1982) Temperature dependences of viscosity of glass-forming substances at constant fictive temperatures. Journal of Non-Crystalline Solids, v. 52, p. 105-114. Mittlefehldt, D. W. (1994) ALH 84001, a cumulate orthopyroxenite member of the Martian meteorite clan. Meteoritics and Planetary Science, v. 29, p. 214-221. Rekhson, S. M., Heyes, D. M., Montrose, C. J., and Litovitz, T. A. (1980) Comparison of viscoelastic behavior of glass with a Lennard-Jones model system. Journal of Non-Crystalline Solids, v. 38-39, p. 403-408. Taylor, G. J. (2000) Liquid Water on Mars: The Story from Meteorites. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/May00/wetMars.html. Taylor, G. J. (1996) Life on Mars? Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/Oct96/LifeonMars.html. Treiman, A. H. (1998) The history of Allan Hills 84001 revised; multiple shock events. Meteoritics and Planetary Science, v. 33, p. 753-764. Turner, G., Knott, S. F., Ash, R. D. and Gilmour, J. D. (1997) Ar-Ar chronology of the Martian meteorite ALH84001: Evidence for the timing of the early bombardment of Mars. Geochimica et Cosmochimica Acta, v. 61, p. 3835-3850.

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PSR Discoveries: Hot Idea: Gusev crater on Mars

http://www.psrd.hawaii.edu/Sept98/GusevMars.html

posted September 24, 1998

For a Cup of Water on Mars: Gusev Crater Written by Linda M.V. Martel Hawai'i Institute of Geophysics and Planetology

Does this martian crater ("cup" in Greek) hold fossils in water-laid sediments?

Though liquid water is not stable on the surface of Mars today, there are hints in the martian landscape that water once flowed there, eroding valleys and depositing sediments. Understanding where water is, or was, on Mars is a crucial step in looking for life on this smaller, cooler neighbor of Earth. Satellite images, beginning in the early 1970s with Mariner 9 up to the current Mars Global Surveyor, have given us increasingly detailed looks at the surface of Mars, including those intriguing channels that resemble dry river valleys. Warmer temperatures and a higher surface pressure once made it possible for liquid water to exist on the surface of Mars. When the water existed and where it went are just two of the questions being studied today. A larger question has to do with martian life. If the wetter, early environment on Mars supported life, then where are the most appropriate places to look for evidence of life? The answer seems to be channels and ancient lake beds. Nathalie Cabrol and colleagues at NASA Ames Research Center, the Vernadksy Institute in Moscow, and Arizona State University recently published their study of a valley and impact crater on Mars which together had a prolonged history of water-related activity. The researchers established a sequence of events for the Ma'adim Vallis/Gusev crater area that included flowing water, ponding, and sedimentation over a period of a couple of billion years. This history makes Gusev crater a prominent depositional site and, as we'll consider, a key location for future biological explorations on Mars. Reference: Cabrol, N. A., E. A. Grin, R. Landheim, R. O. Kuzmin, R. Greeley, 1998, Duration of the Ma'adim Vallis/Gusev Crater Hydrogeologic System, Mars. Icarus, v. 133, p. 98-108.

Carved and jarred surface

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Ma'adim Vallis, in the martian southern cratered uplands, is one of the largest valley networks on the planet. First classified by Sharp and Malin (1975) as a runoff channel, it has long, branched, adjoining channels along its upstream reaches. Flowing northward toward the Elysium Basin region, Ma'adim Vallis is 15 to 20 kilometers wide and extends some 900 kilometers, passing through 160-km-diameter Gusev crater at 14.7oS, 184.5oW. Younger, smaller impact craters on the southern and northern rims of Gusev opened the pathway for water from Ma'adim Vallis.

Cabrol and her colleagues used images from the Viking Orbiter mission to make a geologic map of Ma'adim Vallis and Gusev crater. Through a careful process of distinguishing between different textures, forms, and layers, they were able to relate the surface materials to different geologic processes, namely erosion and deposition by flowing water or (to a lesser extent) wind, and impact cratering.

Unit Descriptions

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The researchers defined five different surfaces: image of unit

name of unit

description and interpretation

Ma'adim Vallis floor

Smooth surface, with some terraced layers. Interpreted to be deposits on the bottom of a river channel.

Ma'adim Vallis terraces and deltaic deposits inside Gusev crater

Grouped together because of similar elevations and surface texture. The multi-leveled terraces are interpreted to be evidence of more than one episode of erosion and carving of the valley.

Gusev floor

Smooth surface interpreted to be deposits of fine-grained sediments transported into the crater by water from Ma'adim Vallis.

Gusev rim

From the crests of the hilly rim to landslide deposits around the base, this unit is related to the actual formation of the crater by an impact event.

Smooth and bright surface at the junction of Ma'adim Vallis and the southern rim of Gusev crater, interpreted to be windblown material. windblown deposits This smooth and bright veneer is probably the youngest material deposited into Gusev crater.

Counting the Craters

After defining the surface units, Cabrol and co-workers used a crater-statistics technique to determine the relative ages of the units and to place the geologic processes into a relative time sequence. In general, older surfaces have more craters simply due to their longer exposure time. They counted the number of craters of three sizes (2-, 5-, and 16-kilometer diameters) in each of the five surface units. These crater populations were used to estimate absolute ages of the surfaces according to Tanaka's model, a technique developed by Ken Tanaka of the U. S. Geological Survey, Flagstaff, as shown in the table below.

From oldest to youngest, Cabrol and her colleagues found the sequence of surface units to be:

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PSR Discoveries: Hot Idea: Gusev crater on Mars

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Unit

Possible age in billions of years

Gusev rim

4

Ma'adim terraces

about 3.5

Gusev floor

about 1.8

Ma'adim floor

about 1.8

windblown deposits

from 2 until now.

The crater statistics show about the same age for the floors of Ma'adim Vallis and Gusev crater, thus constraining the age of the last sedimentation event. Cabrol and colleagues use these statistics to confirm the idea that Gusev was a natural ponding zone and collector basin for Ma'adim Vallis sediment, which could be hundreds of meters thick.

Could there be evidence for life in Gusev's cup of water?

What does the geologic history of Gusev crater and Ma'adim Vallis mean for biological explorations? Biologists, paleontologists, and geologists are studying the current and early environments of Mars to determine if conditions could have supported life. They are, in particular, looking for clues to past or present water. Gusev crater is intriguing because of its long history as a depositional site for water and sediments from Ma'adim Vallis. There is the fascinating possibility that the water-laid sediments in Gusev crater could contain fossils. In 1995, Gusev crater was included in NASA's report, "An Exobiological Strategy for Mars Exploration" (NASA Publication SP-530) as a priority site for future biological exploration. Images taken in 1998 by the Mars Orbiter Camera (MOC) on Mars Global Surveyor show the martian surface in greater detail than previously achieved by the Viking Orbiters. The MOC image below, shown at 40% of the original size, has a resolution of 18.3 meters (60 feet) per pixel.

MOC image link When the Mars Global Surveyor spacecraft attains its Mapping Orbit in early 1999, the MOC will take images

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at an unprecedented scale of approximately 1.5 meters per pixel. Taking close looks at Gusev crater and its surroundings can give us additional information to determine geological details, such as the sources of the sediments, and verify the locations of the ancient pond shorelines. Quality images and geological studies, like those conducted by Nathalie Cabrol and colleagues, will help future mission planners select landing sites for robotic rock collectors and, ultimately, piloted explorations to Mars. It's only a matter of time before we know for sure if the ancient channels and lake beds on Mars hold evidence of past life.

An Exobiological Strategy for Mars Exploration Baker, V. R., 1982, The Channels of Mars. University of Texas Press, 198 p. Cabrol, N. A., E. A. Grin, R. Landheim, R. O. Kuzmin, R. Greeley, 1998, Duration of the Ma'adim Vallis/Gusev Crater Hydrogeologic System, Mars. Icarus, v. 133, p. 98-108. Carr, M. H., 1996, Water on Mars. Oxford University Press, 229 p. Center for Mars Exploration at NASA Ames Space Science Division. Kargel, J. S. and R. G Strom, 1996, Global Climatic Change on Mars Scientific American. Mariner 9 Mission to Mars Mars Global Surveyor MOC high-resolution images MOC image of Gusev crater. Sharp, R. P., and M. C. Malin, 1975, Channels on Mars. Geol. Soc. Am. Bull., v. 86, p. 593-609. Tanaka, K. L., D. H. Scott, R. Greeley, 1992, Global Stratigraphy, in Mars (H.H. Klieffer, B. M. Jakosky, C. W. Snyder, M. S. Matthews, Eds.), p. 345-382, Univ. of Arizona Press, Tucson. Viking Mission to Mars

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PSR Discoveries: Crystal structures or fossils in martian meteorite?

http://www.psrd.hawaii.edu/Dec97/LifeonMarsUpdate2.html

posted December 18, 1997

Fossils in Martian Meteorite: Real or Imagined? Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology

In August, 1996, David McKay and his colleagues at the Johnson Space Center, Stanford University, and McGill University reported evidence for fossil life in a meteorite from Mars. Since then, scientists from around the world have been testing this idea, which in science generally means they are trying to prove the idea wrong. If an idea stands up to such scrutiny, it is on pretty solid ground. The latest round of testing comes from John Bradley (MVA, Inc. and Georgia Tech), Ralph Harvey (Case Western Reserve), and Harry Y. McSween (Univ. of Tennessee). They suggest in a short report in Nature that most of the tiny, tubular, segmented objects that resemble tiny fossils ("nanofossils") described by McKay and his associates are actually thin fractures parallel to atomic planes in minerals in the martian meteorite. The features have been modified, Bradley and coworkers say, by the way samples are prepared for study, producing the apparent segmentations. Dave McKay and his collaborators agree that some structures in the rock are mineral features, but say others are not, and that the sample preparation techniques they used do not produce pronounced segmentation. It appears that a nanofossil is in the eye of the beholder. References: Bradley, J. P., Harvey, R. P., and McSween, H. Y., Jr., 1997, No "nanofossils" in martian meteorite. Nature, v. 390, p. 454. McKay, D. S., Gibson, E., Jr., Thomas-Keprta, K., and Vali, H., 1997, Reply. Nature, v. 390, p. 455.

The Original Observations

The original paper by Dave McKay and his cohorts cited several lines of evidence that hinted at past life in martian meteorite ALH84001, as summarized in the first edition of PSRD. One of them dealt with observations of fossil-like objects in the rock. Though much smaller, the wormy structures resemble ancient microfossils found on Earth. Because of their small size, they are called "nanofossils." in ALH84001

in Earth rocks (from J. W. Schopf)

in ALH84001

Click each image to enlarge it.

The photo on the left, which has appeared in numerous publications, shows a small, tubular, segmented object in ALH84001. Note the resemblance to microfossils found in ancient rocks on Earth (center), although the terrestrial fossils are much larger. In some cases, many of the fossil-like objects in ALH84001 are found in one place, as shown in the

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photo on the right. The curious alignment bothered some scientists, making them wonder if the structures might be related to the crystal structure of the underlying rocky material.

Crystal structures or fossils?

The main argument at first centered on the small sizes of the nanofossils. Some scientists wondered if an organism could actually be that small and still function. Others argued that such small organisms are found in the fossil record on Earth, though some questioned those observations. The viability of such small creatures is still debated, including by Bradley and his colleagues in their communication to Nature, though McKay and friends cite occurrences of bacteria as small as 70 nanometers. The prime tool for examining rocks at high magnification is the electron microscope. There are two basic types. One is called a scanning electron microscope (SEM), which sprays a highly-focused beam of electrons in a grid pattern on the sample, which can be a little chunk of a rock or a polished piece. An SEM is shown in the photograph below. Samples are placed in the sample chamber, which is then evacuated by powerful vacuum pumps. The tall column labeled "electron gun" is also evacuated. It produces electrons near the top and accelerates them down toward the sample. Electromagnets focus the beam to a very fine point (about 1 nanometer in diameter), while other magnets cause it to move in a grid pattern. The monitors allow scientists to view the images produced by detectors that see electrons that bounce off the sample or others (called secondary electrons) that are produced in it. The monitor on the far left provides informtion about the chemical composition of the mineral in the sample. The other type is called a transmission electron microscope (TEM), which transmits electrons through very thin (less than one micrometer) slices of rock. TEMs can also move the electron beam in a grid pattern to produce an image like a picture. An interesting problem in looking at rocks with electron microscopes is that the instruments are so good that you can see features you might not expect to see. Crystallographic features not visible in conventional light microscopes are evident, and some types of fractures can appear curved or worm like. Even more confusing, because the samples need to be coated with some kind of conducting material to carry off the electrical charge deposited by the electrons, scientists have to worry about introducing artifacts when they apply the coating to their samples. Small crystallographic features and artifacts from sample preparation are what Bradley and coworkers argue are the so-called nanofossils observed by McKay and his team.

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Pyroxene Crystal Surface Click on an image to enlarge it.

Carbonate Crystal Surface Click on an image to enlarge it.

Scanning electron microscope images obtained on samples of ALH84001 in John Bradley's lab. Photos on the right are higher magnification pictures of regions in photos on the left. The top pair show the surface of a pyroxene crystal; the bottom pair show a carbonate crystal.

Bradley and his team suggest that the nanofossils in ALH84001 are not fossils at all, but are instead small-scale features of crystals (which Bradley calls "cleavage lamellae") of pyroxene and carbonate in the rock. Basically, these are parallel planes in mineral crystals that leave steps on broken surfaces. Many of the elongated microscopic features in the rock are lined up on the surfaces of crystals and seem to emerge from inside the grains. Bradley says that this is evidence for a crystallographic origin, not fossils deposited on the surface. Some of these features are even curved like the nanofossils reported by McKay. In their reply to Bradley and coworkers, McKay and company agree that there are numerous crystallographic features on crystal surfaces in ALH84001, and they illustrate this with some photographs like those above. They also argue that some features may have been caused by weathering, producing clay minerals. No matter how the linear features formed, McKay and associates agree that such features are not of biological origin. On the other hand, the McKay group points out that there are other curved and more isolated structures besides those shown by Bradley, and it is those structures that the group claims are biological. They also argue that the structures they call nanofossils are much larger, up to 750 nanometers (0.75 micrometers) long, than those shown by Bradley (see the photographs above). John Bradley counters in a communication to PSRD that many of the elongate structures in their images are just as long, and also point out that some of the previous evidence cited by McKay and colleagues, such as the herd of nanomaggots, are the same size as most of the crystallographic features observed. McKay notes that objects like those in the nanomaggot herd are more ovoid than elongate, and thus less likely to be mineralogic features.

Pesty Problems with Sample Preparation

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This is a typical apparatus used for depositing conducting coatings on samples. Samples are placed in the glass bell jar (vacuum chamber) and the jar is pumped down to a low pressure. Rods of carbon or wires of metal such as gold, are heated by electric currents so hot that they begin to evaporate. The vapor lands everywhere, some forming a thin coating on the sample.

Everyone who works with electron microscopes worries about introducing artifacts during sample preparation, and the smaller the feature you examine, the more troublesome the problem becomes. The samples are coated with a conducting substance such as carbon, gold, and palladium. These coatings are ultrathin, typically between 2 and 20 nanometers thick. As thin as that is, it is significant compared with the scale of the features suspected to be nanofossils. McKay and coworkers show an interesting example of this effect. uncoated

coated

The left image, shown by McKay and coworkers in their paper, shows the surface of an uncoated carbonate area in ALH84001. (Surfaces can be examined without coating by using low electron voltages, but the images are not as sharp so the samples are usually coated.) In between the elongated objects the surface is wispy and composed of many tiny grains. The photograph on the right shows the same surface as the left, now coated with a layer of carbon 10 nanometers thick. The elongated objects are unchanged, though easier to see, but the wispy material is no longer visible and the surface looks smooth. This shows that coating samples can hide some features. Another problem is that coating might lead to creating features that were not present in the rock. This is what Bradley and his colleagues suggest has happened to produce the apparently segmented objects in ALH84001. They argue that thick coatings of gold or a mixture of gold and palladium cause the formation of segmented surfaces, and the thicker the coating, the more pronounced the segments. Thus, they argue that the obvious segmented nature of the feature shown in the first photograph in this article is not a real feature of the rock, but formed when the sample was coated.

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McKay and associates disagree with this interpretation as they have observed segmented structures on samples prepared in several different ways. For example, the photograph on the right shows elongated, S-shaped objects with segments, though the segments are not as clear as they are in photos in McKay's original paper. This sample was prepared by making a replicate of the surface by placing a fluid, plastic substance on an uncoated, freshly broken surface of ALH84001, which they believe preserved the original features in the sample. This plastic was removed after it had hardened, then examined in a transmission electron microscope. The objects have a segmented structure, perhaps representing divisions inside the cells, which McKay and colleagues claim is good evidence that the objects are real fossils. John Bradley suggested to PSRD that even the replication technique may introduce artifacts. The technique is well established, but is usually used to replicate objects on the micrometer scale or larger. The problem with replication at the nanometer scale is that the structure of the replicating material itself contributes to the textures you observe. In addition, when the material is removed from the sample, it might form stretch marks or cracks, perhaps producing a bogus segmented structure. Bradley concludes that the most effective way to identify features in rocks as nanofossils is to use transmission electron microscopy, which allows you to see interior cell walls and other diagnostic features.

An Unsolved Problem

It appears that one person's nanofossil is another person's artifact. In spite of this apparent complete disagreement, however, this latest debate has helped clarify some issues. For example, it is agreed that some of the features in ALH84001 that vaguely resemble nanofossils are actually structural features of the minerals, weathering products, or both. It shows that sample preparation is worth an even more thorough look, so that we can know which features might be artifacts and which are probably not. Undoubtedly, the continued debate will lead to new tests of all the evidence for life in ALH84001. The debate over nanofossils is part of a larger debate on the temperature of formation of the carbonates in the meteorite, which is argued even more vigorously than is the reality of the nanofossils! Other PSRD articles that have dealt with the temperature of origin of the carbonates are listed below. Whatever the outcome of all the arguments and detailed study of the martian meteorites, it will show us how to search for life in samples returned from the surface of Mars, which is planned tentatively for the year 2005.

David S. McKay and others, 1996, Search for Past Life on Mars: Possible Relic Biogenic Activity in Martian Meteortie ALH84001, Science, vol. 273, p. 924-930. G. Jeffrey Taylor, "Life on Mars? The Evidence and the Debate." PSR Discoveries. Oct 1996. G. Jeffrey Taylor, "Life on Mars -- The Debate Continues." PSR Discoveries. March 1997. E. R. D. Scott, "Shocked Carbonates may Spell N-o L-i-f-e in Martian Meteorite ALH84001." PSR Discoveries. May 1997. G. Jeffrey Taylor, "Low-temperature Origin of Carbonates Consistent with Life in ALH84001." PSR Discoveries. May 1997. Mars links, including past and future mission descriptions from the National Space Science Data Center (NSSDC).

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PSRD: Hafnium, Tungsten, and the Differentiation of the Moon and Mars

http://www.psrd.hawaii.edu/Nov03/Hf-W.html

posted November 28, 2003

Hafnium, Tungsten, and the Differentiation of the Moon and Mars --- Experiments help us understand the timing of core formation and nature of initial melting in the Moon and Mars.

Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology

Measurements of the isotopic composition of tungsten (W) show that lunar samples and Martian meteorites have an excess of

tungsten-182. This was produced by the decay of hafnium-182 (Hf-182), an isotope with a half-life of only 9 million years. Because tungsten dissolves enthusiastically in metallic iron and hafnium does not, it is possible to use the abundance of W-182 in rocks formed by melting of the silicate mantle as an indicator of the timing of core formation. However, the concentrations of Hf and W in rocky material can be affected by melting and crystallization, so we also need to know how each element concentrates in common minerals in the mantles of the Moon and Mars. The behavior of Hf has been studied experimentally, but this is not true of W. Kevin Righter (Johnson Space Center) and Charles (Chip) Shearer (University of New Mexico) have filled this knowledge void by determining how W partitions between olivine, high- and low-calcium pyroxene, plagioclase feldspar, and garnet. The new data allowed Righter and Shearer to reexamine available measurements of the isotopic composition of W in lunar samples and Martian meteorites. Their analysis suggests that the lunar magma ocean, a huge magma system that surrounded the Moon when it formed, solidified in less than 30 million years. This is shorter than many theoretical calculations suggest. Pathfinder data and chemical data from Martian meteorites suggest that the core of Mars makes up about 20% of the planet. Core formation and subsequent melting of a region of the mantle containing garnet and high-calcium pyroxene took place less than 20-30 million years after the formation of the first solids in the solar system. This type of research shows the importance of measurements of isotopic compositions of radioactive elements or their decay products and laboratory experiments on the geochemical behavior of those elements. References: Righter, K. and Shearer, C. K. (2003) Magmatic fractionation of Hf and W: Constraints on the timing of core formation and differentiation in the Moon and Mars. Geochimica et Cosmochimica Acta, v. 67, p. 2497-2507. Shearer, C. K. and Righter, K. (2003) Behavior of tungsten and hafnium in silicates: A crystal chemical basis for understanding the early evolution of the terrestrial planets. Geophysical Research Letters, v. 30, doi: 10.1029/2002GL015523.

Short-lived Isotopes and Core Formation

Cosmochemists routinely use numerous isotopic systems to determine the ages of solar system materials. The isotopes include both

those with long half lives, such as rubidium-87 ( 87Rb, half life of 48.8 billion years), which decays to strontium-87 ( 87Sr), and those that have half lives that are so short that the radioactive isotope no longer exists. The hafnium-182 ( 182Hf)-tungsten-182 (182W) pair is one of those. 182Hf has a half life on only 9 million years. This means that essentially all of it decayed to 182W in about six half lives, only 55 million years. While this might seem like a long time, it happened 4.5 billion years ago. The decay of a short-lived isotope gives us a way of precisely dating events that happened long ago. Cosmochemists can determine age differences of only a million years in rocks billions of years old. What we are dating depends on the elements involved. In some cases we date the time a solid formed in the solar nebula, the cloud of gas and dust surrounding the nascent Sun. Others might date the time a lava flow erupted onto a planet's surface. Still others can date the time the mantle of a planet first melted. The Hf-W system dates the time metal segregated from rocky materials to form a core and provides a quantitative look at the nature of silicate melting in the newly-formed mantle. The diagram below summarizes the way Hf and W behave as a planet melts early in its history.

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This diagram shows the consequences of formation of a metallic core and subsequent silicate melting on the ratio of 182W to 183W. (This ratio is expressed in the parameter epsilon-tungsten, which normalizes the ratio to that of the primitive chondrite meteorites. In chondrites, therefore, epsilon-tungsten is zero. epsilon-tungsten is also zero in the bulk silicate Earth (crust and mantle combined). Core formation causes the W to concentrate in the metallic core, leaving Hf (and a little W) in the rocky mantle. The core has no 182 Hf in it to make 182 W, so it never reaches the amount in primitive chondrites. This causes epsilon-tungsten to be less than 0. After metal segregates to form a core, the Hf/W ratio is high in the rocky material left behind. This allows lots of 182 Hf to decay to 182 W, leading to epsilon-tungsten larger than zero. There are complicating factors that give cosmochemists the opportunity of delving more deeply into the processes involved in the initial differentiation of the planets. One of those is the behavior of the Hf and W left behind in the rocky portion. Shearer and Righter did experiments to quantify the behavior of W during the melting of silicates (rocks).

Primitive chondrite meteorites and the rocky portion of the Earth form the standard for comparison of the ratio of 182W to 183W. The latter isotope was not formed by the decay of 182Hf and behaves geochemically just like W, so 183W makes a convenient reference. To compare easily, Der-Chuen Lee and Alex Halliday (Institute for Isotope and Mineral Resources in Zurich, Switzerland), pioneers in the analysis of tungsten isotopes, defined a parameter called epsilon-tungsten, or ε w. This simply compares the ratio of 182W to 183W in a sample to the ratio in the bulk silicate Earth multiplied by 10,000. This is expressed mathematically by the equation shown in the diagram above. If the tungsten in a chondrite remains in chemical equilibrium, the 182W/183W will remain chondritic because of tungsten exchange among minerals, so epsilon-tungsten will be zero. However, the situation changes dramatically if metallic melting removes iron to form a dense, metallic core. The tungsten concentrates in the metal, leaving behind a metal-free rock that now has a higher Hf/W ratio. The metal would have virtually no Hf, so no additional 182W would be produced. This leads to a low value of 182W/183W in the metallic core. In turn, this causes epsilon-tungsten to be less than zero. Iron meteorites, which formed in the cores of asteroids that melted, have epsilon-tungsten less than zero.

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Iron meteorites like the one shown here formed when chondrites melted. The metal, because it was denser than the surrounding silicates, sank to form the core of a tiny planet in the asteroid belt. Because the hafnium stayed behind, there was no source of 182W, giving a value of epsilon-tungsten less than 0.

The silicate left behind after metal sank to form a core in an asteroid or planet would still have some 182Hf, unless core formation took so long (more than about six 182Hf half lives, or 55 million years) than none remained. This radioactive 182Hf would continue to decay, producing 182W. Because of this production and the high amount of Hf compared to W, the ratio of 182W to 183W becomes larger than chondrites have, so epsilon-tungsten is greater than 0. How much greater depends on several factors such as when the core formed and the resulting Hf/W ratio. The Hf/W ratio can be affected by further melting in the rocky portion of a planet. The ratio in a magma or in the unmelted residue can vary substantially depending on what minerals are present. Cosmochemists, particularly Gordon McKay (Johnson Space Center), have done experiments to map out the way hafnium behaves during melting, but few experiments have been done to understand how tungsten behaves. Shearer and Righter decided to remedy this knowledge gap.

Measuring Element Behavior

Laboratory experiments to determine the geochemical behavior of elements have helped us understand a great deal about the

geochemical evolution of the planets. Trace elements are particularly useful because they behave in predictable ways, once some experimental data are available. Cosmochemists can calculate the way elements separate from each other. In the case of radioactive elements, cosmochemists can even date the time when the elements separated. Kevin Righter took the lead in doing experiments on how tungsten partitions into magma and different minerals as the rock inside a planet melts. The expected minerals to be involved in melting inside planets are olivine, low-calcium pyroxene, high-calcium pyroxene, plagioclase feldspar, and garnet. The idea is to concoct a synthetic rock of the right composition, then heat it until it is partially melted. To be applicable to the interiors of planets, the experiments must be done at elevated pressure. Righter used pressures ranging from 1 bar (atmospheric pressure at the surface of the Earth) to 100 kilobars (a kilobar is 1000 times the atmospheric pressure). Temperatures ranged from 1150 Celsius to 1850 Celsius. Part of the reason for doing the experiments at a range of pressures is that some minerals, such as garnet, form only at high pressure. The high-pressure apparatus, called a multi-anvil, uses octahedral devices that are squeezed between tungsten carbide cubes. (The tungsten carbide is used because of its strength, not because Righter and Shearer were studying the behavior of tungsten. It does not interfere with the experiments on the behavior of tungsten inside the ceramic octahedral.) Some experiments at the low end of the pressure range were done with another type of pressure apparatus called a piston-cylinder device.

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To determine the behavior of tungsten at high temperature and pressure, Kevin Righter used high-pressure apparatus at the University of Arizona. The samples to be studied were placed in the center of ceramic octahedral (white) and then squeezed between eight truncated tungsten carbide cubes (grayish, metallic). (The tungsten does not interfere with the experiments on the behavior of tungsten inside the ceramic octahedral.) Depending on the sizes of the octahedrons and the truncations on the tungsten carbide cubes, pressures between 3 and 30 billion pascals (30,000 to 300,000 times atmospheric pressure) can be obtained. This corresponds to depths of ~100 to 1000 km on Earth and 300 to 3000 km on Mars.

The assembled stack of eight cubes and the sample octahedron is placed in the center of this cylindrical module, between six wedges that are connected to a hydraulic system. The blue module is called a "multi-anvil module."

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http://www.psrd.hawaii.edu/Nov03/Hf-W.html

The multi-anvil module is placed into a 1000 ton press in the laboratory of Professor Mike Drake (right) at the University of Arizona. Drake is a pioneer in determining partitioning behavior of trace elements from high-temperature and high-pressure experiments. Nancy Chabot (left) and Chris Capobianco (right), then at the University of Arizona, observe a sample being pressurized.

Righter and Shearer used a range of starting compositions so they could be sure of getting all the minerals of interest to crystallize and to study how the tungsten partitioning varies with mineral composition. They also added about 0.5 wt% WO 3 (and in some cases HfO2) to be sure there would be enough in all minerals and the magma to measure accurately. To ensure that conditions were not oxidizing (appropriate for planetary interiors), Righter encapsulated the samples in graphite-lined platinum tubing. The graphite forces the system to maintain oxidizing conditions at the right level. It also prevents or at least minimizes loss of iron from the sample to the platinum capsules. (Iron loss would change the chemical properties of the minerals being studied.) Righter and Shearer measured the chemical compositions of the experimental products on an electron microprobe, a standard device that does chemical analyses on tiny spots (only 1 micrometer across). Chip Shearer measured the tungsten abundances using the ion microprobe at the University of New Mexico. Ion microprobes are very complicated instruments used to measure both isotopes and trace elements. Shearer had to work out assorted interferences from other elements and create a set of standards. Ion microprobe analyses are rarely routine!

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http://www.psrd.hawaii.edu/Nov03/Hf-W.html

The concentrations of tungsten were measured with the ion microprobe at the University of New Mexico. This complicated instrument consists of an ion source, mass spectrometer, and associated electronics and vacuum system. Shown operating the ion probe is Justin Hagerty, a graduate student at UNM.

The partitioning behavior is described by a parameter called the partition coefficient. This is simply the ratio of the concentration of an element in a mineral crystal divided by its concentration in the magma in which the mineral is forming. A partitioning coefficient of 1.0 indicates that the element would have the same concentration in the mineral as in the magma. Partition coefficients less than 1 shows that the element prefers to remain in the magma; those greater than 1 indicate that the element concentrates in crystals in preference to magma. The partition coefficients for tungsten and hafnium determined by Righter, Shearer, and others appear in the table below. Mineral

Distribution Coefficient for W

Distribution Coefficient for Hf

Olivine

0.02 to 0.07

0.07

Low-Ca pyroxene

0.007 to 0.025

0.08

High-Ca pyroxene

0.04 to 0.19

0.28 to 0.56

Plagioclase

0.003 to 0.03

--

Garnet

0.007 to 0.028

0.20 to 0.52

All the distribution coefficients for tungsten are very low, showing that tungsten concentrates in magma. (However, if metallic iron is present, the tungsten readily concentrates in it.) Distribution coefficients for hafnium are low for olivine and low-Ca pyroxene. This means that those minerals will not cause drastic separation of tungsten from hafnium. This is far from true for high-calcium pyroxene and garnet, however. Tungsten avoids them, but hafnium, though still preferring a magma, concentrates relatively easily into those minerals. This means that when a rocky body melts after metal has segregated to a core, the presence of high-calcium pyroxene and garnet will lower the Hf/W ratio in the magma and increase it in the unmelted part of the rock. Assuming this happens when 182Hf is still present, it will lead to lower epsilon-tungsten in the magma and higher epsilon-tungsten in the solids. This allows cosmochemists to figure out the depth and timing of melting early in a planet's history.

Making Sense of Tungsten Isotopes

Righter and Shearer applied their data to published isotopic measurements of lunar basalts and Martian meteorites. The lunar samples contain excesses of 182W; epsilon-tungsten ranges from 1 to about 4.5. To interpret the data, they calculated what the epsilon-tungsten value would be if core formation happened 30, 40 and 50 million years after the formation of the chondrites. They also calculated the effect of different amounts of fractionation of hafnium from tungsten (expressed in the figure below by different Hf/W). High Hf/W

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could have resulted from accumulation of high-calcium pyroxene (cpx in the graph below) and ilmenite (iron-titanium oxide). Lunar scientists believe that high-titanium lunar basalts formed from mantle rocks that formed by accumulated of high-calcium pyroxene and ilmenite in the lunar magma ocean. This would explain the high epsilon-tungsten value one of them has. Low-titanium basalts formed from mantle rocks that accumulate olivine and low-calcium pyroxene, leading to low Hf/W, and low epsilon-tungsten. In both cases, the calculations fit the data best if the core formed no later than 30 million years after the formation of the solar system. This means that materials in the solar nebula accreted into larger and larger objects, forming the Earth and other planets, and then the Moon (probably by a giant impact onto the Earth). As the Moon assembled in Earth orbit, it melted, forming the magma ocean. The magma ocean cooled and crystallized, forming a feldspar-rich crust and dense cumulates that later remelted to form the basalts that make up the maria. And all that happened in 30 million years or less!

The left side of this diagram shows how epsilon-tungsten would vary with time in the Moon, depending on when the core of the moon-forming impactor formed and on the Hf/W ratio after formation. Because Hf concentrates preferentially over W in ilmenite (ilm) and high-calcium pyroxene (cpx), rocks rich in those minerals will have high Hf/W. Such rocks could form late during the solidification of the lunar magma ocean. When rocks rich in ilmenite and high-calcium pyroxene melted, the magma would also have high epsilon-tungsten. Parts of the Moon that contained only olivine and low-calcium pyroxene would have lower Hf/W and lower epsilon-tungsten. The right side of the diagram shows epsilon-tungsten values for four lunar mare basalts. The data suggest that the magma ocean had to have solidified within 30 million years of the formation of the first solids in the solar nebula. Otherwise, epsilon-tungsten in low-titanium basalts such as 15555 would be lower. Formation of the high-titanium basalts seems to have involved mantle rocks rich in ilmenite and high-calcium pyroxene.

Martian meteorites also have positive values of epsilon-tungsten (see graph below). Righter and Shearer again calculated how epsilon-tungsten would change with time depending on when the core formed and the Hf/W in the mantle. In contrast to the smaller Moon, Mars probably has garnet in it (see cross section of Martian interior). Since garnet is one of the minerals that might greatly affect the ratio of Hf/W, hence the value of epsilon-tungsten, Righter and Shearer looked at several cases. A deep part of the mantle rich in garnet would have very large epsilon-tungsten, exceeding 20. The Martian meteorites do not have values anywhere near that. This does not mean that a mantle like did not form on Mars, only that the Martian meteorites did not form from that type of mantle rock. An alternative is that the small excess of 182W (small positive epsilon-tungsten) is inherited from a shallow mantle source that contained high-calcium pyroxene and perhaps garnet. For example, a mantle containing consisting of 10% high-calcium pyroxene and 10% garnet (with the remainder being olivine and orthopyroxene) would have a Hf/W of 17. Such a mantle could produce epsilon-tungsten values close to those in Martian meteorites (shown on the right). They also calculate that a garnet-free mantle could have Hf/W of about 7 and would also be consistent with the epsilon-tungsten measured for Martian meteorites. Whatever the details, it appears that the Martian mantle and core could have formed rapidly, easily within 20-30 million years of the beginning of the formation of the solar system.

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The left side of this diagram shows how epsilon-tungsten would vary with time in Mars, assuming core formation at 15 million years after the beginning of the solar system. The calculated curves are for different Hf/W ratios in the mantle, depending on mineralogy (see diagram below). Because Hf concentrates preferentially over W in garnet, a deep mantle rich in garnet would have high Hf/W. The resulting epsilon-tungsten values would be much larger than in Martian meteorites (right side of diagram). Smaller values like those in Martian meteorites can be obtained if the meteorites come from shallow parts of the mantle that contain less garnet (Hf/W of 7 to 17). Even with the low Hf/W of 7, epsilon-tungsten values like those in the meteorites are obtained within 20-30 million years of the origin of the solar system.

Experiments by Constance M. Bertka and Yingwei Fei (Geophysical Laboratory of the Carnegie Institution of Washington) give us one possible picture of the interior of Mars. In this picture, the uppermost mantle of Mars consists of olivine and pyroxene, with a small amount of garnet (shaded green). However, at a depth of about 1100 km, the olivine begins to convert to a more dense form, called gamma-spinel, without changing its chemical composition. The conversion is complete by 1300 km. Along with the conversion of olivine to a spinel crystal structure, garnet and pyroxene convert to a mineral called majorite, which has a crystal structure like garnet, but is close to pyroxene in chemical composition (shaded yellow). At higher pressures, hence deeper, there is a relatively abrupt transition at 1850 km (shaded black) to a mixture of perovskite (itself a mixture chemically of MgSiO3 and FeSiO3) and magnesiowustite (a mixture of FeO and MgO). The metallic core (shaded gray) begins at about 2000 km depth and continues to the center at a depth of 3390 km. Magma produced by melting in any of the silicate regions above the metallic core will result in different ratios of Hf to W because the two elements partition differently into different minerals.

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y

The Knowledge Matrix

The study of extinct isotopic signatures like the Hf-W system used by Righter and Shearer required several important advances and tools. One is the development of instrumental techniques to measure small differences in the isotopic composition of tungsten. The second is development of high-pressure experimental apparatus. The third is the development of the ion microprobe for trace element analyses of tiny grains. Then these techniques had to be used to measure the isotopic compositions in lunar samples and meteorites, and to determine the geochemical behavior of hafnium and tungsten. Once all that was done, cosmochemists (in this case Kevin Righter and Chip Shearer) synthesized the data into a new understanding of the timing and processes involved as the planets formed and began to evolve geologically. Most cosmochemical research involves such interdisciplinary advances in instrumentation, data, and understanding.

Halliday, Alex N. (2000) Hf-W chronometry and inner solar system accretion rates. Space Science Reviews, v. 92, p. 355-370. Righter, K. and Shearer, C. K. (2003) Magmatic fractionation of Hf and W: Constraints on the timing of core formation and differentiation in the Moon and Mars. Geochimica et Cosmochimica Acta, v. 67, p. 2497-2507. Shearer, C. K. and Righter, K. (2003) Behavior of tungsten and hafnium in silicates: A crystal chemical basis for understanding the early evolution of the terrestrial planets. Geophysical Research Letters, v. 30, doi: 10.1029/2002GL015523.

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9 of 9

PSR Discoveries: Rootless cones on Mars

http://www.psrd.hawaii.edu/June01/lavaIceMars.html

posted June 26, 2001

If Lava Mingled with Ground Ice on Mars Written by Linda M.V. Martel Hawai'i Institute of Geophysics and Planetology

Clusters of small cones on the lava plains of Mars have caught the attention of planetary geologists for years for a simple and compelling reason: ground ice. These cones look like volcanic rootless cones found on Earth where hot lava flows over wet surfaces such as marshes, shallow lakes or shallow aquifers. Steam explosions fragment the lava into small pieces that fall into cone-shaped debris piles. Peter Lanagan, Alfred McEwen, Laszlo Keszthelyi (University of Arizona), and Thorvaldur Thordarson (University of Hawai`i) recently identified groups of cones in the equatorial region of Mars using new high-resolution Mars Orbiter Camera (MOC) images. They report that the Martian cones have the same appearance, size, and geologic setting as rootless cones found in Iceland. If the Martian and terrestrial cones formed in the same way, then the Martian cones mark places where ground ice or groundwater existed at the time the lavas surged across the surface, estimated to be less than 10 million years ago, and where ground ice may still be today. Reference: Lanagan, P.D., A. S. McEwen, L. P. Keszthelyi, and T. Thordarson (2001) Rootless cones on Mars indicating the presence of shallow equatorial ground ice in recent times, Geophysical Research Letters, vol. 28, p. 2365-2368.

Location and Description of Martian Cones

Cone-shaped structures on the Martian volcanic plains were first identified and interpreted as rootless cones in the 1970s with Viking imagery. Their occurrences were reported in Chryse Planitia, Deuteronilus Mensae, Acidalia Planitia, Isidis Planitia, and Elysium Planitia. Using the higher-resolution MOC images, Lanagan and colleagues identified cones in the Cerberus plains, Marte Valles, and Amazonis Planitia. The cones appear to be superimposed on the surface of low volcanic plains near recognized outflow channels. Using crater counts and other geologic evidence, William Hartmann (Planetary Science Institute) determined the lava flows may be as young as 10 million years.

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This regional topographic map of the area around Cerberus Plains has white contour lines (at 200-meter intervals) superimposed over a Viking image. Low areas have darker shading. White arrows indicate downstream directions of known outflow channels. The small black squares with white dots mark locations of cone clusters identified by Lanagan and colleagues in MOC images.

The cones seen in the MOC images have base diameters ranging in size from 20 meters to 300 meters. Summit craters on the cones have diameters about half as wide as the bases. Martian cones are found in clusters ranging from a few to many hundreds of cones. These dimensions and arrangements are consistent with explosive rootless cones found in Iceland.

Overlapping and clustered cones (a) on Mars and (b) in Iceland. The scale of the Martian and terrestrial cones are comparable.

How Rootless Cones Form

The mingling of lava and water is a violent interaction creating explosions and ejections of hot lava fragments. The 2 of 4

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process begins when lava flows across a wet landscape, perhaps a marsh, shallow lake, shallow aquifer or possibly even icy ground. Although the surface of the lava cools to a hardened crust, molten lava continues to surge through internal tubes. Water that comes into contact with the hot lava becomes heated and flashes to steam. When the steam pressure exceeds the pressure of the lava above it: boom. Repeated explosions build cone-shaped structures of ash and larger pieces of spattery lava. Crowned with summit pits, the cones sit on top of the lava crust. These cones are called "rootless" because they are not fed by a magma source vent lying directly underneath. Lava moving laterally through internal tubes to the explosion site feeds rootless cones. The cones are characteristically distributed in small clusters with no obvious alignment along fissure vents and they do not erupt lavas themselves. Thordarson has found that Icelandic rootless cones are clustered in marshy regions where water-saturated sediments and lava mix thoroughly. By analogy, Lanagan and co-authors say that rootless cones on Mars may require sufficient shallow ground ice to produce water-saturated sediments.

Raudholar Rootless Cone Group at Reykjavik, Iceland. This group is located about 1.5 kilometers from the eastern edge of town (see apartment buildings on the horizon for scale.) The cone group is in the Leitin lava flow, which covered a small shallow lake here. The cluster is located about 27 kilometers from the source vents of the lava.

Implications for Recent Equatorial Ground Ice on Mars

If the cones identified on Mars are evidence of geologically recent shallow ground ice near the equator of Mars then we can begin to make some assumptions about how far down the ice might have been. Based on calculations by C. Allen (Johnson Space Center), lava-water explosions on Mars driven by steam pressure would require the water or ice to be at a depth of no more than half the thickness of the lava flow. Lava flows in the Marte Valles region have been mapped at 10-meters thickness. Therefore, any ground ice available for creating rootless cones on Mars was no deeper than 5 meters below the surface lava. Shallow ground ice present on Mars less than 10 million years ago, the researchers argue, would mean that deposits of shallow ground ice probably persist near the cones today. Lanagan and colleagues list three possible sources for the shallow Martian ground ice. It could be relic ice leftover from the planet's formation. It could be recondensed water vapor from the ground-atmosphere water vapor exchange. It could be recharge from surface flooding events. The researchers favor the third case, citing the proximity of the cones to outflow channels. They note that it is unlikely that relic ground ice has survived for four billion years in equatorial regions of Mars, but perhaps plausible that vapor exchange between the ground and atmosphere, as modeled for Mars by Sarah Fagents and Ronald Greeley (Arizona State University), was sufficient to recharge the ground ice. Confirming the presence of shallow ground ice in the equatorial region of Mars would be an exciting discovery. A satellite launched April 7, 2001 carries the technology to collect the data we need. The Gamma Ray Spectrometer onboard the 2001 Mars Odyssey mission will supply data on the distribution and abundance of chemical elements at or 3 of 4

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near the surface of Mars, and may provide the key to finding water ice in the shallow subsurface. Its expected arrival at Mars is October 24, 2001, with a primary science mission scheduled for January, 2002 through July, 2004. Finding ground ice and understanding its distribution, if it truly exists, may soon be within our reach.

Allen, C. C.(1980) Volcano-Ice Interactions on the Earth and Mars, in Advances in Planetary Geology, NASA TM-81979, p. 161-264. Fagents, S. A. and R. Greeley (2000) Formation of pseudocraters on Earth and Mars (abstract), Volcano-Ice Interactions on Earth and Mars, p.13. Hartmann, W. K. and D. C. Berman (2000) Elysium Planitia lava flows: Crater count chronology and geological implications, Journal of Geophysical Research, vol. 105, p. 15011-15025. Lanagan, P.D., A. S. McEwen, L. P. Keszthelyi, and T. Thordarson (2001) Rootless cones on Mars indicating the presence of shallow equatorial ground ice in recent times, Geophysical Research Letters, vol. 28, p. 2365-2368. Thordarson, T. (2000) Rootless eruptions and cone groups in Iceland: Products of authentic explosive water to magma interactions (abstract), Volcano-Ice Interactions on Earth and Mars, p. 48. 2001 Mars Odyssey.

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4 of 4

PSRD: Cosmochemistry Instrument: LIBS

http://www.psrd.hawaii.edu/Oct06/libs.html

posted October 26, 2006

Written by Linda M. V. Martel Hawai'i Institute of Geophysics and Planetology

In this series of articles, "Instruments of Cosmochemistry," PSRD highlights the essential tools and amazing technology used by talented scientists seeking to unravel how the solar system formed. You will find information on how the instruments work as well as how they are helping new discoveries come to light.

Laboratory Caliber Instrument on a Rover

Laser-induced breakdown spectroscopy (LIBS) is an active remote sensing technique used for the rapid characterization of elemental compositions of materials. Used for years in laboratory and industry applications, it will make its debut performance on rocks and soils on another planetary surface in 2010 as part of the ChemCam instrument package onboard NASA's Mars Science Laboratory (MSL) rover scheduled for a 2009 launch to Mars. A combined Raman-LIBS is also planned to be part of the Pasteur instrument payload on the ExoMars rover mission planned by the European Space Agency for a 2011 launch. In preparation for use on Mars, a team of scientists at Los Alamos National Laboratory, Roger Wiens, Justin Thompson, James Barefield, David Vaniman, Sam Clegg, and colleague Horton Newsom (Institute of Meteoritics at the University of New Mexico) have tested the LIBS technique on two Martian meteorites and a terrestrial analog rock. Their work confirms that LIBS is capable of determining even subtle differences in rock types from a stand-off distance of 5.4 meters. This high-quality remote sensing on the surface of Mars is exactly what's needed to push the state-of-the-art of cosmochemical investigations as we prepare for follow-up Mars sample return missions. Reference: Thompson, J. R., R. C. Wiens, J. E. Barefield, D. T. Vaniman, H. E. Newsom, and S. M. Clegg (2006) Remote Laser-Induced Breakdown Spectroscopy Analyses of Dar al Gani 476 and Zagami Martian Meteorites. Journal of Geophysical Research, v. 111, doi: 1029/2005JE002578,2006.

How LIBS Works

Laser-induced breakdown spectroscopy (LIBS) uses a high power pulsed laser, focused on the target, to provide more than a megawatt of power on a small spot less than a millimeter diameter for a few billionths of a second. The target rock can be up to 13 meters away from the instrument (otherwise known as the stand-off distance). Each laser pulse vaporizes thin layers of the target rock--a process known as laser ablation--producing a hot spark or plasma. This supersonically expanding plasma glows with electronically excited ions, atoms, and small molecules from the target rock (see image below.) 1 of 8

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Click on the movie frame to view a QuickTime movie in a new window. This picture shows the glowing LIBS plasma produced in air during a laboratory test where the laser was five meters away from the rock. The high-temperature ablated material breaks down into electronically excited atoms and ions, giving off light when they decay back to lower energy levels. The light emitted by the plasma can be collected and analyzed through spectrometers to resolve the characteristic emission lines of the elements that are present in the target rock. Original source: http://libs.lanl.gov/LIBS_movies.html. Note that the apparent wandering of the plasma position on the rock is due to motion of the rock during the test. There is no positional instability of the laser relative to the spark size.

The plasma light is collected by a reflecting telescope and directed through a fiber-optic cable to spectrometers, which resolve and measure the elemental emission lines in the plasma spectrum. In a typical analysis, the spectra from multiple pulses (for example 75 to 100 pulses) are averaged for greater statistical accuracy into one final spectrum for the analysis spot. The LIBS technique yields detailed, quantitative information on compositions of the elements (high and low atomic numbers), including some minor and trace elements, that are present in the target rock. This information is obtained very quickly, within minutes, and will allow scientists to identify rocks on the surface of Mars that are of greatest interest and may be chosen for further investigation by instruments that require physical contact or for collection. Why LIBS is an outstanding tool for planetary surface analyses: no sample preparation is required operates at a stand-off distance (typically 2-13 meters), which permits remote analysis of inaccessible rocks (perhaps up on cliff) the laser removes dust from target surfaces, again without the need to drive to and touch the surface repetitive laser pulses on the same analysis spot permits ablation down through weathering rinds to measure composition through depth profiling and examine the pristine rock chemistry simultaneous multi-element detection (major, minor, trace elements) rapid analysis (a typical analysis sequence is six minutes) good detection sensitivity; 10 ppm detection limits for some elements laser requires only an average of 3 Watts of power during the several minutes of instrument operation time (For more details see the ChemCam Fact Sheet produced by Los Alamos National Laboratory. Link opens in a new window.)

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This view of a cliff of layered rocks at Victoria crater was made by the panoramic camera on NASA's Mars Exploration Rover Opportunity in September, 2006. The cliff is about six meters (20 feet) tall. In the future, scientists will be able to use LIBS to analyze the rocks up on the cliff that are out of reach of the rover. Theoretically, LIBS could produce elemental profiles of the entire wall of the outcrop. [Click the image for higher resolution versions.]

LIBS Cosmochemical Applications to Mars

The ability to identify and quantify the elemental compositions of rocks and soils on Mars is of paramount importance to understanding important issues of the planet's formation and alteration. LIBS spectra will play a vital role in allowing scientists to address these key issues: igneous processes and what they tell us about planetary differentiation and evolution of magma compositions through time sedimentary processes and what they tell us about the interactions between rock and water or atmosphere hydrothermal and weathering processes that have modified (or are currently altering) the Martian crust and what they tell us about the history of water on Mars movement and deposition of materials climate and habitability of Mars

In preparation for use on the Martian surface, the research team at Los Alamos and University of New Mexico tested the LIBS technique in their laboratory on natural rock samples under simulated Martian conditions. The ability of LIBS to distinguish between rocks of widely differing compositions is well known. But what Thompson, Wiens, and colleagues wanted to test specifically was the ability to remotely distinguish a range of igneous compositions on Mars. So, they chose to study two basaltic Martian meteorites with slightly different compositions and textures and an andesite rock powder standard. The team analyzed Dar al Gani 476 and Zagami, two basaltic shergottite meteorites. Both samples were in the form of sawn slabs, shown below.

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Dar al Gani 476 is a basaltic shergottite with olivine and pyroxene phenocrysts up to 5 millimeters in size set in a fine-grained groundmass of average grain size of 0.13 millimeters. This sample is roughly 1 centimeter across.

Zagami is a basaltic shergottite described as a composite of up to three and minor related lithologies shock-melted glass. Overall this rock is finer grained than DaG 476. This sample is roughly 1.5 centimeters across.

The LIBS analyses were conducted under simulated Martian conditions, with the Martian meteorites and andesite sample each placed in a vacuum chamber maintained at a static pressure of ~7 Torr CO2. Each analysis spot was shot 100 times with a pulsed Nd:YAG laser resulting in ablation pits that ranged in diameter from 0.4 to 0.5 millimeter. Fourteen analysis spots were recorded on DaG 476 (five are shown in the figure below) and nine spots were used on Zagami (also shown below). Thompson and his coauthors chose a larger number of analyses on DaG 476 to try to compensate for the rock's larger grain sizes.

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These pictures show some of the LIBS analysis spots or pits (~400 micrometers diameter) on slabs of DaG 476 (top left) and Zagami (top right). Each pit was produced by 100 pulses of the laser beam. The red arrow from spot 2 on DaG 476 points to a magnified back-scattered electron image of the pit. Each LIBS spectrum used by the research team was produced by averaging 100 laser pulses.

A typical LIBS spectrum for an analysis spot on DaG 476 collected by the research team is shown below.

Major elemental emission lines in the LIBS spectrum for DaG 476 are labeled in this plot. The Roman numeral in parentheses refers to the ionization state of the atom. (I) is an excited neutral atom, (II) is a singly-ionized atom. LIBS spectra obtained for Martian meteorite Zagami and the andesite rock powder standard 5 of 8

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(JA-2) are similar at the level of detail shown. The spectrum was collected from a stand-off distance of 5.4 meters.

The diagrams below show how well the LIBS results compare with the average of published chemical analyses (averages of several analyses) of the two Martian meteorites and the andesite. In general, the LIBS results are within 10% of the oxide compositions reported in the literature.

Elemental compositions obtained by the research team using LIBS (shown as weight percent on the x-axis) are compared to literature whole-rock compositions (y-axis) of DaG 476, Zagami, and andesite JA-2. An exact match in oxide composition between the LIBS tests and published laboratory results would plot on the straight line. The oxides plotted here are SiO 2, FeO, CaO, MgO, TiO2, Al2O3, and Na2O. The lower plot is an enlargement of the left half of the top plot. For clarity, we have labeled the CaO and MgO data points, which are discussed below. Click here to see a table of values that were used to create these plots (table will open in a new window.)

In their study, Wiens, Thompson, and colleagues demonstrated that the LIBS technique is capable of determining subtle differences in rock types from a remote distance of 5.4 meters. The differences between LIBS and values determined by traditional laboratory techniques differed by less than 12% (relative) for most of the major elements 6 of 8

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(most were better than 6%). The agreement for titanium and sodium were not as good, but their concentrations are low (less than 1 wt%), hence the emission signals are low. The comparison with the andesite was particularly good with all elements agreeing to within 5% (relative), except for titanium. Even for titanium the difference between the LIBS and accepted concentration for the andesite was only about 8%. The andesite was a powdered sample, hence homogeneous on the scale of the LIBS laser pit. In contrast, the Martian meteorites were slabs of rock. Getting a good average analysis in this case requires numerous analytical points because sometimes the laser zaps only one mineral, other times a mixture, and so on. The most important result of this test of LIBS analytical capability is that it is possible to distinguish between the two Martian meteorites. This is particularly noteworthy for MgO and CaO, as shown in the diagram above. DaG 476 and Zagami are clearly different in composition. This shows that LIBS will be capable of distinguishing rock types as the rover journeys across the Martian surface.

What's Planned for LIBS on Mars

LIBS instruments are planned for NASA's Mars Science Laboratory rover (scheduled to launch in the fall of 2009) and ESA's ExoMars rover (slated for launch in mid-2011).

Click on the images for more information about these future missions.

Mars Science Laboratory rover is designed to operate for a full Martian year, which is almost two Earth years. LIBS, as part of the ChemCam instrument package, is expected to make thousands of measurements to help scientists characterize the geology of the landing region, help analyze surface ices or salts or evaporite minerals or rocks and soils that have been altered by water, help identify possible organic materials if they exist or ever existed, and help check for toxic materials. A typical analysis sequence for LIBS will begin when the science team identifies a target rock and commands ChemCam to fire a burst of up to 75 laser pulses at a ≤ 1 millimeter spot on the target. The rover's onboard spectrometers will determine the elemental compositions of the ablated plasma. Acquiring a LIBS spectrum for an analysis spot is expected to take six minutes. This is very rapid compared with previous techniques that have taken up to three Martian sols for analogous dust-free analyses that required contact with the target rock. Like all scientific instruments, LIBS has to be calibrated using standards of known composition. A challenge on Mars faced by the ChemCam science team is that they will not have a suite of calibration standards that they can expose at the same conditions (distance, etc.) Roger Wiens and team are currently working on determining how the calibration curves vary with instrument-to-sample distance. They will have some calibration standards on the rover, which will be at a close range of 1.4 meters. Over time, the team will also be able to cross-calibrate between LIBS and the other ChemCam instruments APXS and CheMin. [See this ChemCam web page for additional information on the science instruments.] These aspects should allow LIBS to be an extremely useful quantitative geochemical and geological mapping tool. 7 of 8

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The research team at Los Alamos and University of New Mexico who tested LIBS on the Martian meteorites and terrestrial analog rock, as well as other teams of cosmochemists who are working to better calibrate LIBS for the entire variety of rocks it will encounter on Mars, are eagerly anticipating what they'll see in the glowing LIBS light. LINKS OPEN IN A NEW WINDOW.

ChemCam LIBS instrument description from Jet Propulsion Laboratory. ExoMars mission homepage from the European Space Agency. LIBS planetary science applications website from Los Alamos National Laboratory. Mars Science Laboratory rover homepage from Jet Propulsion Laboratory. Thompson, J. R., R. C. Wiens, J. E. Barefield, D. T. Vaniman, H. E. Newsom, and S. M. Clegg (2006) Remote Laser-Induced Breakdown Spectroscopy Analyses of Dar al Gani 476 and Zagami Martian Meteorites. Journal of Geophysical Research, v. 111, doi: 1029/2005JE002578,2006. [ About PSRD | Archive | Search | Subscribe ] [ Glossary | General Resources | Comments | Top of page ] 2006 [email protected] main URL is http://www.psrd.hawaii.edu/

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PSR Discoveries: Hot Idea: The Martian Interior

http://www.psrd.hawaii.edu/Aug97/InsideMars.html

posted August 22, 1997

The Martian Interior Written by G. Jeffrey Taylor

(NASA photo)

Hawai'i Institute of Geophysics and Planetology

The Viking missions in the mid-1970s and the recent Mars Pathfinder mission have shown us what the surface of Mars is like at three places. Their instruments have revealed the chemical composition of the surface and individual rocks. Exciting as those missions were, they literally only scratched the surface. Sojourner's Alpha Proton X-Ray Spectrometer (APXS) examined rocks and soil on Mars. Here the rover is atop Mermaid Dune. (Image MRPS #82938, taken on Mars Sol 30, NASA; see enlargement.)

One of the most important things to learn about the planets is the nature of their interiors because much of the story of their formation and geological evolution is recorded in the chemical composition and minerals inside the rocky planets. Estimates of the chemical composition of the interior of Mars have been made on the basis of the compositions of martian meteorites, some chemical reasoning, and judicious assumptions. However, until recently, the actual minerals present at different depths in the interior could only be guessed because no comprehensive experiments had been conducted at the high pressures and temperatures appropriate to the interior of Mars. Those experiments have now been done by Drs. Constance M. Bertka and Yingwei Fei of the Geophysical Laboratory of the Carnegie Institution of Washington. Although application of their experiments to Mars still requires some assumptions about how temperature varies with depth and the composition and size of the metallic core at its center, Bertka and Fei suggest that the martian mantle has two main layers, one extending from a depth of about 50 km (at the base of the crust) to around 1100 km, a second from 1100 km to about 1800 km, and a third thin layer occuping a zone about 100-200 km thick above the metallic core. Reference: Bertka, Constance M. and Yingwei Fei, 1997, Mineralogy of the martian interior up to core-mantle boundary pressures. Journal of Geophysical Research, vol. 102, p. 5251-5264.

Using

to probe its interior

We have samples of igneous rocks from Mars, in the form of the SNC meteorites. They were blasted off by impacts millions of years ago, and eventually made their way to Earth to fall as meteorites. (The best evidence that they come from Mars is that two of them contain trapped gas with a composition matching that of the atmosphere of Mars as measured by the Viking landers; see Meteorites from Mars at the Johnson Space Center.) Of the 12 martian meteorites, all of which solidified near the martian surface by crystallization from a cooling magma, the most useful probes of the interior are the shergottites (the S in SNC). How do chemical

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analyses of martian meteorite samples tell us about the planet's innards? Magmas originate inside planets, but their compositions are different from the rock that melts to produce them. Some elements are readily incorporated into magmas, while others stay behind in the unmelted portions of the mantle. As the magmas ooze their way to a planet's surface, other things can happen to them, such as partial crystallization in magma chambers beneath volcanoes, further changing their compositions. In spite of this complicated chemical behavior, the abundances of elements in magmas contain a wealth of information about the composition of the region of the mantle in which they formed. The key to decoding this information lies in knowing that not all elements are greatly affected by melting and crystallization. Many elements that concentrate in the magma behave chemically alike, so the ratio of one to another does not change. It's as if magmas retain a detailed memory of their birth. But the decorder key works only if you can determine the composition of a magma. The martian meteorites allow us to do just that...to analyze the compositions of martian magmas. This basaltic shergottite meteorite from Mars, named Zagami, has a mass of 18 kilograms and was observed to fall in Nigeria, Africa and recovered in 1962 (see enlargement). Samples like Zagami are investigated in great detail in geochemical laboratories around the world.

Chemical Make-up

While several approaches have been tried to solve the puzzle of the martian interior, the one generally considered to be the best guess was made in the mid-1980s by G. Driebus and H. Wanke of the Max Planck Institute in Mainz, Germany. Their basic approach was to measure the abundances of a few key compounds, such as FeO (iron oxide), in shergottite samples, and then make use of averages and element ratios to estimate the abundances of the elements in Mars itself. Iron is an extremely important element because it occurs as FeO in planetary mantles and as metallic Fe in cores. (It also occurs in more oxidized form (Fe2O3) on the surfaces of planets, appearing reddish. However, the amount of Fe2O3 is too small to worry about when computing the composition of an entire planet.) FeO does not concentrate strongly into magmas when they form, nor does it stay behind preferentially in the unmelted solids. Its abundance in a magma, therefore, is not much different from its abundance in a planet's mantle. The average FeO in five shergottites is 18.9 wt.%, implying that the martian mantle contains roughly that amount. Driebus and Wanke were not satisfied with that rough estimate. They also noted that the MnO (manganese oxide) abundance in shergottites averaged 0.48 wt.%, about the same as in carbonaceous chondrites. They used this to make the bold assumption that the mantle of Mars has exactly the same MnO content as carbonaceous chondrites. Then, dividing the ratio of FeO/MnO in the shergottites (39.5) by the same ratio in carbonaceous chondrites (100.6), they estimate that the FeO content of the martian mantle is 0.39 of the FeO content of the chondrites. This translates to an abundance of FeO of 17.9 wt.% in Mars.

This example of a carbonaceous chondrite was collected in the Grosvenor Mountains, Antarctia. It is 6.4 x 3.6 3.3 cm in size and weighs 106.2 grams (see enlargement).

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Why would Driebus and Wanke think that carbonaceous chondrites had anything to do with Mars? Because these dark, carbon-rich, grundgy meteorites have chemical compositions very similar to the Sun (except for the most gaseous elements such as hydrogen and helium), they are thought to represent the average stuff that makes up the Solar System, including the planets. The chondrites form a useful basis for comparison. By making numerous other comparisons of element abundances, Driebus and Wanke were able to deduce the concentrations of the most abundant elements in Mars, and some less abundant ones too (see table below). A comparison with the composition of the Earth's mantle (which is much better known) shows that Mars contains much more FeO than does the Earth. Compositions of the rocky portions of Mars and Earth (wt.%) Mars

Earth

SiO2

44.4

45.1

TiO2

0.1

0.2

Al2O3

3.0

4.0

Cr2O3

0.8

0.5

MgO

30.2

38.3

FeO

17.9

7.8

MnO

0.5

0.1

CaO

2.4

3.5

Na2O

0.5

0.3

K2O

0.04

0.03

Compound

The composition and size of the metallic core is also important. The larger it is, the smaller the depth and lower the pressure at the boundary between the mantle and core. Driebus and Wanke estimated the composition of the martian core by assuming that the total composition of Mars is the same as carbonaceous chondrites and that all the iron and nickel not present in the mantle are in the core as metallic Fe. They also assumed that the core contains all the sulfur in the planet. This gives a composition of the martian core of 77.8 wt.% iron, 7.6 wt.% nickel, and 14.2 wt.% sulfur. Although the high FeO composition of the martian mantle is generally accepted by planetary scientists, not all agree. In particular, S. Ghosal, R. O. Sack, and M. E. Lipschultz of Purdue Univeristy, and M. S. Ghiorso of the University of Washington in Seattle, argue on the basis of mineral compositions in martian meteorites that conditions in the martian mantle were much less oxidizing than generally assumed. This would produce more metallic iron in the core, and less oxidized iron (FeO) in the mantle. Using a computer program called MELTS that simulates melting and crystallization of magmas, they estimate that the martian mantle contains only 6.8 wt.% FeO, even lower than that of the Earth. Their work, to be published in Contributions to Mineralogy and 3 of 6

PSR Discoveries: Hot Idea: The Martian Interior

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Petrology, will be hotly debated when it appears in print. PSR Discoveries will follow the debate as it unfolds. So, now we have estimates of the chemical composition of the Martian interior. But what about the minerals present? This is where Bertka and Fei enter the story.

Producing the pressures and temperatures inside Mars

To know the physical state of materials at the high pressures and temperatures existing inside planets, you must simulate those conditions in the laboratory. This requires some highly specialized equipment that can compress and heat samples, and then allow them to be recovered for study. Bertka (in photo, see enlargement) and Fei used a device called a multi-anvil press. It is made of a sturdy housing that contains six removable wedges. The wedges transmit forces from a hydraulic piston to the interior of the device. Inside there is a cube made of eight separate tungsten carbide cubes, each with a corner cut off. The truncated corners meet in the center, forming a region shaped like an octahedron. This is where the high pressures are attained during an experiment. A small furnace (less than 1 cm in diameter) containing the sample is placed in the octahedral area inside the multi-anvil press. The samples Bertka and Fei used were a mixture of oxide powders with overall compositions identical to the calculated composition of the martian mantle. Once assembled, the furnace could be controlled to give the desired temperature, and the pressure controlled by the force applied by the hydraulics. The experiments lasted between 3 and 48 hours at temperatures of 1100oC to 1765oC and pressures ranging from 20,000 to 235,000 times the pressure at the surface of the Earth. These conditions correspond to depths of 200 to 2000 km inside Mars. (See note about pressure.) After each experiment, the samples were removed from the now-crushed furnaces, mounted in epoxy, and polished. These mounts were then examined in an electron microprobe, a device that can measure chemical compositions in small spots, allowing identification of the minerals present. The results show that the minerals change as pressure increases, even though the chemical composition of the material remains the same. Not surprisingly, at higher pressures, minerals with higher densities form. The mineral olivine, no longer stable at the higher pressure (20 GPa), converts to a crystal form called gamma-spinel. The pyroxenes and garnet convert to majorite. These mineralogies are illustrated in the false-color images below.

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Electron microprobe image of an experimental product colored according to chemical composition. Pressure was 2 GPa. Temperature was 1100oC. The experimental run lasted for 48 hours.

In this case, the pressure was 20 GPa. Temperature was 1750oC. The experimental run lasted for 3 hours.

The Interior of Mars

The experiments of Bertka and Fei give us one possible picture of the interior of Mars (see enlargement). In this picture, the uppermost mantle of Mars consists of olivine and pyroxene, with a small amount of garnet (shaded green). These are fairly common minerals on Earth, the other planets, the Moon, and asteroids. However, at a depth of about 1100 km, the olivine begins to convert to a more dense form, called gamma-spinel, without changing its chemical composition. The conversion is complete by 1300 km. Along with the conversion of olivine to a spinel crystal structure, garnet and pyroxene convert to a mineral called majorite, which has a crystal structure like garnet, but is close to pyroxene in chemical composition (shaded yellow). At higher pressures, hence deeper, there is a relatively abrupt transition at 1850 km (shaded black) to a mixture of perovskite (itself a mixture chemically of MgSiO3 and FeSiO3) and magnesiowustite (a mixture of FeO and MgO). The metallic core (shaded gray) begins at about 2000 km depth and continues to the center at a depth of 3390 km.

One difficulty in arriving at the above picture from the high pressure experiment is a great uncertainty in the rate at which the temperature increases with increasing depth inside Mars. Bertka and Fei used a temperature profile that assumes that the core is still molten. Different profiles have been assumed by other investigators, which can lead to slight differences in the inferred depths at which major mineralogical changes take place. Of all the features in Bertka and Fei's picture, the presence of the thin, lowermost mantle layer is most uncertain. If the core is cooler than Bertka and Fei assume, then the layer may not exist. It's presence may affect the formation of zones of mantle upwelling, called plumes, that could have produced widespread volcanism on

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Mars. The dynamics of plume formation is dependent on the nature of the lower mantle, so this is an important matter to settle.

Testing the Model

How can we test the picture Bertka and Fei have produced of the Martian interior? The most direct test would be to install a network of seismometers on the surface of Mars, to record marsquakes. This would show how seismic wave velocities vary with depth, from which the density and temperature profiles could be established. The data would also allow measurement of the size of the core and thickness of the crust. Mars missions, including Pathfinder, will also improve estimates of a parameter called the moment of inertia factor, which provides information about the distribution of mass inside a planet. This parameter is not known with certainty now, so different estimates for the composition of the interior of Mars cannot be tested. Once we have seismic data and a better determination of the moment of inertia factor, it will be possible to decide if the high-FeO mantle favored by Dreibus and Wanke is correct, or whether the low-FeO mantle composition recently proposed by Ghosal and coworkers is more reasonable. A global network of seismometers would also show how the mantle of Mars varies horizontally. When this is compared to Bertka and Fei's internal structure, any variations might be due to motions in the mantle, perhaps associated with incipient plate tectonics, or rising plumes that drove volcanism.

The Mars Pathfinder Sojourner rover analyzed the rock called "Yogi" on Mars Sol 6. (Image no. 81314, NASA.) A seismic network around Mars could someday be installed by rovers. Thus, Mars Pathfinder is an important first step in the exploration of the surface and interior of Mars.

Bertka, Constance M., and Yingwei Fei, 1997, Mineralogy of the Martian interior up to core-mantle boundary pressures. Journal of Geophysical Research, vol. 102, p. 5251-5264. Dreibus, G., and H. Wanke, 1985, Mars: A volatile-rich planet. Meteoritics, vol. 20, p. 367-382. Lindstrom, M., "Martian Meteorites."Meteorites from Mars. 1996. Longhi, J., E. Knittle, J. R. Holloway, and H. Wanke, 1992, The bulk composition, mineralogy, and internal structure of Mars. In Mars (eds. H. H. Kieffer, B. M. Jakosky, C. W. Snyder, and M. S. Matthews), p. 185-208. Univ. of Arizona Press, Tucson. McSween, H. Y., Jr., 1994, What we have learned about Mars from SNC meteorites. Meteoritics, vol. 29, p. 757-779.

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PSR Discoveries: Hot Idea: Life on Mars

http://www.psrd.hawaii.edu/Oct96/LifeonMars.html

posted October 18, 1996

Life on Mars? Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology

People have long wondered if life could have existed or even still exits on Mars. The Viking landers in 1976 searched for signs of life in the red soil, but found no clear-cut evidence. Future missions are planned to search other terrains on Mars, such as areas where water must have flowed in rivers and formed lakes that eventually dried up. But the search has already started. A group of investigators at the Johnson Space Center and Stanford University has revealed evidence from an intense, careful study of a meteorite from Mars that tiny bacteria-like creatures may have lived in cracks in the rock. In this first issue of PSR Discoveries, we describe evidence the researchers have assembled, and present some of the nonbiological alternatives other scientists have proposed. We intend to follow the debate as it unfolds during the coming months or, perhaps, years. Begin your discovery here. Investigate the type of evidence of most interest to you or simply go through the list in order.

Hot Idea Contents Meteorite from the Ancient Crust of Mars. ALH 84001 originated as a slowly-cooled igneous rock in the Martian crust, was excavated by an impact, altered by fluids, and finally sent to Earth by another impact. The Evidence and the Debate. The NASA-Stanford group cites four lines of evidence for fossil life in ALH 84001, all of which are found associated with unusual globules of carbonate minerals: (1) formation before arrival on Earth; (2) concentrations of organic chemicals; (3) tiny grains of iron oxide and iron sulfide; and (4) tubular, fossil-like objects. Alternative interpretations appear in pop-up windows which you open with a click of a button (a JavaScript enhancement). Ancient Hospitable Mars. Before about 3 billion years ago, Mars may have had a more clement climate than now, perhaps allowing life to develop.

The Researchers. Reference: McKay, David S., and others (1996) Search for Past Life on Mars: Possible Relic Biogenic Activity in Martian Meteorite ALH84001, Science, v. 273, p. 924-930.

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PSR Discoveries: Hot Idea: Life on Mars

posted October 18, 1996

Life on Mars? Meteorite From the Ancient Crust of Mars

by G. Jeffrey Taylor

Photograph of a specimen of Allan Hills (ALH) 84001, a 1.9 kilogram (4.2 pound) meteorite found in Antarctica. The little cube in the picture is 1 cm across. Like most meteorites, it was partly covered with smooth, dark, glassy material, called the fusion crust, which formed when the rock blazed through the Earth's atmosphere. It was found in 1984 during the annual meteorite search in Antarctica. According to geologist Roberta Score, former laboratory manager in the meteorite curatorial facility and the explorer who actually found the meteorite, the rock looked greenish inside as it lay on the Antarctic ice. In the laboratory, however, it looks gray. (NASA photo.)

View of a thin slice of ALH 84001 in a microscope reveals large crystals (up to 6 mm long) of orthopyroxene (a silicate mineral containing iron and magnesium) and a small grain near the top of the photo of plagioclase feldspar (sodium-calcium alumino-silicate), rendered glassy by shock waves. Orthopyroxene makes up about 95% of the rock, and the large size of the crystals suggests that the rock crystallized in a slowly-cooling magma body inside the Martian crust. The crystals contain numerous cracks and are separated by crushed zones of much smaller crystals. These zones probably formed when high-pressure shock waves, generated by an impact, crushed portions of the large crystals. Crushed zones and other cracks in the rock contain the carbonate globules that have the features ascribed to biological processes. (Photo courtesy of David Mittlefehldt, Lockheed Engineering and Science Company.)

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PSR Discoveries: Hot Idea: Life on Mars

The meteorite is decorated with globules of carbonate minerals that seem to occur along cracks in the rock. These globules have a somewhat orange color and are small, only 0.1 millimeter across. As discussed in The Evidence and the Debate, a big discussion centers on the origin of the globules, especially whether they formed from very hot fluids (more than 650 degrees Celsius) or cooler ones (between 0 and 80 degrees Celsius) . Life would not have survived high temperatures. (NASA photo.)

When viewed in an electron microscope, it is obvious that the carbonate globules are complicated. This photograph is a colorized image of the intensity of electrons bounced back from a polished surface of a sample of ALH 84001. The colors represent different minerals. Green is orthopyroxene (the silicate with iron and magnesium), blue is glassy plagioclase feldspar, and the various shades of red and orange are carbonate minerals with a range in chemical composition. (Photo courtesy of Ralph Harvey, Case Western Reserve University.)

Studies of ALH 84001 have revealed the basic outline of the rock's history. It formed about 4.5 billion years ago in a relatively large magma body inside the crust of Mars. Its high abundance of one mineral (orthopyroxene) indicates that this mineral must have accumulated in the magma, probably near the bottom of the magma body, eventually forming the original igneous rock with large crystals of orthopyroxene. (Graphics by Brooks Bays, PSR Discoveries graphic artist.)

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PSR Discoveries: Hot Idea: Life on Mars

An impact blasted ALH 84001 4.0 billion years ago, ripping it from its deep location and probably placing it nearer to the surface in a pile of debris. The shock waves deformed the pyroxene crystals and converted the feldspar to glass. This event also heated the rock, allowing Ar gas to escape and resetting the potassium-argon clock, which allows scientists to determine the age of the impact. On the basis of the elemental compositions of the carbonate minerals, Ralph Harvey (Case Western Reserve University) and Harry Y. McSween (University of Tennessee) have proposed that the rock was 650-700 degrees Celsius after the impact and hot fluids rich in carbon dioxide circulated through the crater, depositing the carbonate globules along cracks. (Graphics by Brooks Bays, PSR Discoveries graphic artist.)

In contrast to Harvey and McSween, most investigators, such as Allan Treiman of the Lunar and Planetary Institute and others at the Johnson Space Center in Houston and the Open University in England, believe that mineral compositions and the abundances of the isotopes of carbon and oxygen in the globules imply that the carbonates were deposited by relatively cool (no more than 80 degrees Celsius) flowing water enriched in carbon dioxide, after the rock had been deformed by impact. Determining the age of the carbonate globules is extremely difficult. Estimates range from 1.4 to 3.6 billion years. The age is not known accurately enough to link the formation of the carbonates to the 4.0 impact event, to the relatively wet era on Mars between 3 and 4 billion years ago, or to any time before it was blasted off Mars and sent our way. (Graphics by Brooks Bays, PSR Discoveries graphic artist.)

Scientists in Switzerland, Japan, and the U.S. (Arizona, and California) have measured the time ALH 84001 was exposed to cosmic rays in space. This actually dates the time the meteoroid containing the rock was smaller than a few meters across; the interiors of larger objects are shielded from radiation. This time is between 16 and 17 million years ago, and may indicate when it was lifted off Mars by an impact as depicted in this artist's rendition. It could have been liberated earlier, however, as a large object, and the 16 to 17 million years simply dates a recent breakup of the object as it wandered in space. (Graphics by Brooks Bays, PSR Discoveries graphic artist.)

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It is easy to determine how long a meteorite has been on Earth if it was seen to fall. Fortunately, we can also determine the residence time of other meteorites by determining the extent to which radioactive isotopes (produced by cosmic rays) have decayed. Useful isotopes for this purpose are carbon-14 and aluminum-26. Measurements done on ALH 84001 by scientists in Arizona show that the meteorite fell about 13,000 years ago. It was eventually spotted in 1984 by Roberta Score, and identified as a Martian meteorite in 1994 by one of Roberta's colleagues, Dave Mittlefehldt. Now ALH 84001 is the focus of intense scientific scrutiny because of the possibility that the carbonate globules were formed in part by biological activity of ancient Martian life forms.

Go to the Evidence and the Debate. Return to Life on Mars? [ About PSRD | Archive | Search | Subscribe ] [ Glossary | General Resources | Comments | Top of page ]

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PSRD Hot Idea: Life on Mars

posted October 18, 1996

Life on Mars? The Evidence and the Debate

by G. Jeffrey Taylor

Researchers at the Johnson Space Center and Stanford University have outlined four main lines of evidence that point to biological activity in Allan Hills (ALH) 84001 and, hence, that life once existed on Mars. They have carefully pointed out alternatives to each piece of the story, but they argue that taken as a whole, the best explanation for ALL the features associated with the orange carbonate globules in ALH 84001 is that tiny organisms lived on the surfaces of cracks. Other scientists have devised even more alternatives than David McKay and his colleagues had considered, and a spirited debate is beginning. PSR Discoveries presents the evidence assembled by McKay and coworkers, along with the alternatives that have been discussed so far. Whatever the outcome of the debate, it seems certain that it will spark a great deal of research and we will end up knowing more about the geologic, and possibly biologic, history of Mars.

The Evidence and the Debate Contents Carbonate globules formed on Mars.

Carbonate globules formed from liquid water.

Polycyclic aromatic hydrocarbons.

Tiny grains of magnetite and iron sulfide.

Microfossils.

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PSRD Hot Idea: Life on Mars

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October, 1996 by G. Jeffrey Taylor

There is little doubt that a group of meteorites come from Mars. The evidence is described clearly in a report on the web site of the meteorite curators at the Johnson Space Center (see link on our title page). The ratios of the isotopes of oxygen indicate that ALH 84001 is also a member of that group, though it is much older than the others (over 4 billion years vs 1.3 billion years or less). The meteorite also contains trapped gases like those in the Martian atmosphere. It seems highly likely that ALH 84001 comes from Mars, and there is not much debate about this point. This is an artist's rendition of the impact that liberated ALH 84001 from the martian surface. (Graphic by Brooks Bays, PSR Discoveries graphic artist.)

There is also little dispute that the carbonate nodules formed before the meteorite arrived on Earth 13,000 years ago. As seen in this photograph, which has been colorize to highlight compositional distinctions, the prominent chemical zoning pattern is offset in some carbonate globules, undoubtedly because of an impact event that took place before arrival on Earth. An arrow points to the offset, shown prominently by the white bands. The formation age of the carbonates also indicates a pre-terrestrial origin for them. Although the age of the globules is highly uncertain, there is no question that they formed at least a billion years ago, long before the meteorite landed in Antarctica. (Image based on photograph by David Mittlefehldt.)

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October, 1996

by G. Jeffrey Taylor

Survival of life like that on Earth requires hospitable conditions, most notably water. Data on the isotopic compositions of oxygen in the carbonate globules indicate that the carbonates formed between 0 and 80 degrees Celsius, appropriate for life to flourish. (NASA photo)

Alternative View of Temperature

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PSR Discoveries: Hot Idea: Life on Mars

October, 1996 by G. Jeffrey Taylor

Organisms are made of complicated hydrocarbons (compounds made mostly of hydrogen and carbon), so their presence should be marked by high concentrations of hydrocarbons produced when the organisms decayed. One group of hydrocarbons produced by decomposition of ancient organisms on Earth are called polycyclic aromatic hydrocarbons. These are certainly aromatic: they stink! The simplest one is benzene, depicted here. The corners of the hexagonal structure are occupied by carbon atoms, and a hydrogen atom is bonded to each carbon. The structure of benzene is usually drawn as a hexagon with a circle in the center. The circle represents six electrons in the molecule that are not associated with specific carbon atoms, but are spread out above and below the plane containing the carbon atoms. (Graphic by Brooks Bays, PSR Discoveries graphic artist.)

More complicated aromatic hydrocarbons consist of benzene molecules linked together, such as phenanthrene, shown here. When two or more benzenes are joined the compounds are called "polycyclic aromatic hydrocarbons," or PAHs for short. A number of PAHs were detected in ALH 84001. Researchers at Stanford University, working with colleagues at the Johnson Space Center, have shown that the PAHs in ALH 84001 are not contaminants from the laboratory or Antarctica. PAHs are produced by decay of organic materials; for example, PAHs are abundant in coal deposits. Their presence in ALH 84001 suggest to the Stanford-NASA team that organisms were present. The researchers acknowledge that PAHs are also present in carbon-rich meteorites and in interplanetary and interstellar dust, in which PAHs formed by nonbiological chemical processes, but show that the PAHs in ALH 84001 are different from those in other meteorites, except for a type called "CM carbonaceous chondrites." CM chondrites contain clay-like minerals, organic compounds, magnetite, and iron sulfides. Astronomical observations of asteroids suggest that many asteroids may be like CM carbonaceous chondrites. (Graphic by Brooks Bays, PSR Discoveries graphic artist.)

Alternative Views of Nonbiological Hydrocarbons

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October, 1996 by G. Jeffrey Taylor

McKay and co-workers have identified very small grains of magnetite (iron oxide) and two types of iron sulfide. These have similar sizes and shapes as magnetite and iron sulfide grains formed by bacteria on Earth. This photo shows an iron sulfide grain from the Martian meteorite (left) and a similar grain in a terrestrial bacteria living in the cell of a plant root. (Photo adapted from Science.)

Alternative View of Mineral Formation

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Whether the shapes can be produced by non-biological processes or not, McKay and colleagues argue that the types of minerals present and evidence for some of the carbonate dissolving suggests that biological activity was involved. This photograph shows the distribution of small magnetite (left) and sulfide grains in a carbonate matrix. (Photo adapted from Science.)

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This photograph shows a light band cutting across a carbonate grain. McKay and co-workers suggest that this band was formed by partial dissolution of carbonate. It is in these areas that the magnetite and iron sulfides shown above are found. According to the research team, dissolution of the carbonate required that the water be acidic, but formation of magnetite and iron sulfide from water would have required alkaline (far from the acidity needed to dissolve carbonate), unless bacteria or other microorganisms were involved. The lack of a simple non-biological way to produce the minerals existing together leads them to conclude that magnetite and sulfide formed as the result of biological processes. (Photo adapted from Science.) Alternative View of Mineral Dissolution

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PSR Discoveries: Hot Idea: Life on Mars

October, 1996 by G. Jeffrey Taylor

The most stunning evidence for most of us is the presence of tiny, tube-shaped objects that resemble terrestrial microfossils. The one shown here, photographed with an electron microscope, is about 20 nanometers wide (that's only 0.00002 millimeters) and has segments suggestive of filamentous cyanobacteria. Cyanobacteria used to be called blue green algae. They occur as single-cellular or multicellular (filamentous) forms. Bacteria and cyanobacteria are called prokaryotes, which are organisms whose cells do not have a nucleus. Instead, they have a single strand of DNA, strung in a closed loop. (NASA photo.)

These very tiny fossils were discovered in Western Australia by J. William Schopf (University of California, Los Angeles). A photograph of each specimen appears with an interpretive drawing of the structure. The resemblance to the structures in ALH 84001 is quite striking, though those from the Earth are much larger: the scale bar on the photo is in micrometers, rather than nanometers. (Photo courtesy of J. William Schopf.)

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This image shows a large number of microscopic fossil-like objects on ALH 84001 resembling a herd of nanomaggots. Each one is about 10 nanometers long. (NASA Photo.)

Alternative View of Shapes

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13 of 16

PSR Discoveries: Hot Idea: Life on Mars

posted October 18, 1996

Life on Mars? Ancient Hospitable Mars by G. Jeffrey Taylor All known forms of life on Earth require the presence of liquid water. Mars is an attractive planet to search for extraterrestrial life because its surface contains clear evidence that water flowed across it. There are large channels and valley networks, both of which seem to require large amounts of flowing water. The meteorites from Mars contain hydrated minerals, indicative that water was present in their magmas, hence available to be transferred to the atomsphere to produce a far wetter climate than possessed by present-day Mars. How much warmer and wetter the atomsphere was is not known with certainty, but there certainly was abundant flowing water, especially early in Martian history.

Large channels like this one in Kasei Vallis indicate that water once flowed in prodigious amounts on Mars. However, this does not imply that it had to be incredibly rainy on Mars. In fact, it may not be possible to form such huge floods by rainfall alone. The water more likely emerged from the ground when ice melted rapidly, perhaps because of magmas moving through the crust. The water would end up spurting from the ground, sweeping downhill and eroding the landscape. (23oN, 65oW, NASA photo.)

This photograph of an area near the mouth of Ares Vallis in Chryse Planitia shows the power of the surging water. Flood waters flowing from the bottom to the top of the image were diverted by two craters 8-10 kilometers in diameter. Two streamlined islands were formed. (20oN, 31oW, NASA photo.)

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Valley networks also indicate the presence of liquid water on the surface of Mars. Some may have been formed by groundwater flowing onto the surface, but others resemble typical branching drainage networks on Earth. However, this and other networks on Mars lack the small-scale streams feeding into larger ones. This may indicate that rainfall was not the only process at work to provide the water to carve the valleys. (42oS, 92oW, NASA photo.)

This branching, or dendritic, drainage network in South Yemen was photographed by the Space Shuttle. Note that it is more intricate than the network on Mars, with many smaller streams flowing into larger ones. Calculations suggest that the amount of water required to form channels and valley networks on Mars could have been a few percent of the volume of Earth's oceans, although some estimates place the amount at much less than one percent. On Mars as on Earth, there would have been seas and land masses, not a global ocean.The presence of water on Mars, at times flowing in great rivers and standing in lakes (which were probably frozen on top), makes it promising to search for life on this desert-like, reddish planet. (NASA photo.)

Most of the prominent valley networks occur in the ancient highlands of Mars. This region is characterized by numerous large craters that have been strongly eroded. Since most large craters formed before about 3.8 billion years ago (an age inferred from studies of large craters on the Moon and from lunar samples), erosion rates must have been quite high, certainly much higher than they have been since that time. ALH 84001 is an old rock, formed in the ancient highlands and was involved in a large cratering event 4.0 billion years ago. Conditions in the ancient highlands would have made it likely that the rock was exposed to water, either on the surface or flowing through cracks beneath the ground. (NASA photo.)

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PSR Discoveries: Life on Mars: The Researchers

posted October 18, 1996

Life on Mars? The Researchers List of authors of the Science article: McKay, David S., and others, 1996, Search for Past Life on Mars: Possible Relic Biogenic Activity in Martian Meteorite ALH84001, Science, vol. 273, p. 924-930. David S. McKay Johnson Space Center

Everett K. Gibson, Jr. Johnson Space Center

Kathie L. Thomas-Keprta Lockheed Martin

Hojatollah Vali McGill University

Christopher S. Romanek University of Georgia

Simon J. Clemett Stanford University

Xavier D. F. Chillier Stanford University

Claude R. Maechling Stanford University

Richard N. Zare Stanford University

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PSR Discoveries: Hot Idea Update: Life on Mars?

http://www.psrd.hawaii.edu/Mar97/LifeonMarsUpdate.html

posted March 31, 1997

Life on Mars--The Debate Continues Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology

An exhilarating debate is taking place about whether a team of NASA and Stanford University scientists found evidence for life in a meteorite from Mars, as they claim, or whether their evidence can be attributed to completely nonbiologic processes or to contamination. PSR Discoveries summarized the original evidence for life in a martian meteorite in October, 1996. Since then, many scientists have joined in the discussion and they have obtained new data for the martian meteorite at the center of the debate (ALH 84001, shown here) and other meteorites found in Antarctica. New data involve: additional analysis of hydrocarbons in meteorites, and electron microscope examinations of tiny whisker-like, magnetite grains (Fe3O4) in ALH 84001. Most of the data tends to argue against some of the pieces of evidence that led to the bold interpretation that there were fossils in ALH 84001. Nevertheless, the great debate is far from over! Each piece of evidence will be tested in several ways.

Hydrocarbons: biological or contaminants?

One of the chief lines of evidence used by David McKay and his coworkers at the Johnson Space Center and Stanford University was the presence in ALH 84001 of a distinctive array of polycyclic aromatic hydrocarbons (PAHs). These are smelly, difficult-to-pronounce compounds composed of linked chains of benzene rings. Phenanthrene is shown here as an example (drawn by Brooks Bays, PSR Discoveries graphic artist). McKay and associates argued that the relative abundances of various types of PAHs in ALH 84001 were different from abundances produced by nonbiological processes, such as found in some types of meteorites and interplanetary dust. They concluded that the abundances of PAHs present in the martian meteorite most likely formed by the decomposition of biological compounds.

Several investigators have pointed out that there are other ways of looking at the data. Edward Anders of the University of Chicago argues that all the PAHs observed in ALH 84001 can be made by nonbiological processes, depending on the proportions of carbon, hydrogen, and oxygen, and on the temperature and pressure at which the reactions took place. He also points out that a common type of carbon-bearing meteorite, CM chondrites, contain PAHs very much like those in ALH 84001, and there is no evidence for fossil life in those meteorites. Jeffrey Bell (a colleague of mine here at the University of Hawai'i) also noted the similarity of the PAHs in ALH 84001 and CM chondrites, and even suggested a source for

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abundant CM material on the surface of Mars: chips from the disrupted precursors of Phobos and Deimos, the two satellites of Mars. McKay's colleagues at Stanford, led by Simon Clement and Richard Zare, acknowledge these arguments, but argue that nonbiologic synthesis of PAHs at low temperatures (less than 150 oC) would not produce the distributions observed. On the other hand, if the temperatures did exceed 150 oC, Clement and Zare agree that nonbiologic processes, such as the Fischer-Tropsch reaction, would produce PAH abundances like those measured in ALH 84001. Thus, determining the temperature at which the carbonates formed is very important, and this is one of the most hotly contested points in the debate.

It may not matter whether the PAHs formed at high or low temperatures, if Luann Becker, Daniel Glavin, and Jeffrey Bada of Scripps Institute of Oceanography are right. (Dr. Becker is also with the Space Science Division at NASA Ames Research Center.) They argue that all the PAHs in ALH 84001 are contaminants. Dave McKay and his colleagues had conducted some rigorous tests for contamination, concluding there was none, but Becker and coworkers' work raises suspicions. G. D. McDonald and J. Bada reported last year that they had found amino acids in another martian meteorite, designated EET 79001. This rock is a basaltic lava flow, and quite young, only 180 million years old. (ALH 84001 formed about 4.5 billion years ago.) Amino acids are complex organic compounds made and used by living organisms to make proteins. They are also found in carbonaceous chondrites. However, the amino acids in carbonaceous chondrites are a mixture of left- and right-handed molecules (a reference to the symmetry of the molecules). Organisms on Earth use L-amino acids, the left-handed variety. McDonald and Bada found only L-amino acids in EET 79001, which they interpreted as contaminants. Why not conclude that the L-amino acids were produced on Mars by martian organisms? McDonald and Bada preferred the contamination interpretation because they also found similar amino acids in Antarctic ice. Luann Becker's experiment was to measure the PAHs in EET 79001, and compare them to those in ALH 84001 and Antarctic ice. She found that the PAHs in carbonate minerals in EET 79001 were almost identical to those in ALH 84001. The PAHs measured in the ice were also similar to those in the meteorites, though not identical.

Comparison of PAHs Data from Becker and Bada (Geochemica et Cosmochimica Acta, January, 1996) Compound and atomic mass ALH 84001 EET 79001

Ice

naphthalene (128)

not detected not detected present

fluorene (166)

not detected not detected not detected

phenanthrene or pyrene (178) present

present

present

chrysene (228)

present

present

present

perylene (252)

present

present

present

anthanthrene (276)

not detected present

not detected

anthanthracene (278)

present

present

coronene (300)

not detected not detected present

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The PAHs in both meteorites are associated with carbonate mineral deposits. It is these carbonates in ALH 84001 that contain the evidence for fossil life, so why not simply conclude that Becker and Bada have discovered some evidence for life in a second martian meteorite? They argue that the small carbonate grains adsorb PAHs from any water in which they come in contact. To test that idea, Becker added carbonate to distilled water that contained a known amount of PAHs, let it interact overnight, and then extracted the carbonate for analysis. The results show that virtually all the PAH added to the water had concentrated into the carbonate grains, demonstrating how easily the carbonate soaks up PAHs, whatever the source. From these measurements, Becker and coworkers concluded that the PAHs in both ALH 84001 and EET 79001 were contaminants, though they could not determine if the contamination occurred on Mars or in Antarctica. In either case, they believe that the source of the PAHs is probably material like the CM chondrites and interplanetary dust, which rains down on both Earth and Mars. Of course, water has not been flowing recently on Mars, so if the contamination took place there, it did so at a time when the climate on Mars was much different than it is now. The experiments support Bell's idea that the source of the CM-like PAHs were fragments of a disrupted ancient martian moon. The presence of terrestrial L-amino acids in EET 79001, though, argues that the contamination took place on Earth.

Of course, these arguments could be turned around. The similarity in the PAHs in the two meteorites and the Antarctic ice, and the presence of Earth-like amino acids in EET 79001 might indicate that life on Mars closely resembles life on Earth. And not explained as yet is the unusual isotopic composition of the carbon in ALH 84001, which Ian Wright, Monica Grady, and Colin Pillinger (Open University and Natural History Museum, England) suggest cannot be caused by contamination. The debate about the meaning of the organic analyses of the meteorites will continue, and more studies are planned, including measurements of other organic compounds and their isotopic compositions. Whiskers of magnetite

Carbonate areas in ALH 84001 contain tiny grains of magnetite, Fe3O4, which McKay and coworkers interpreted as having been manufactured by martian microorganisms. Some bacteria on Earth produce such tiny magnetite grains. However, a close look at the magnetites in ALH 84001 by John Bradley (MVA, Inc. and Georgia Institute of Technology), Ralph Harvey (Case Western Reserve University), and Harry Y. McSween (University of Tennessee) suggests that the magnetites formed by nonbiological processes at relatively high temperatures, 500 to 800 oC. Bradley and his colleagues used an analytical scanning transmission electron microscope to examine the magnetites. These high-tech devices allow scientists to study the shapes, crystal structures, and chemical compositions of grains smaller than 1 nanometer (one billionth of a meter). Most of the magnetites in the carbonate areas do not have distinctive diagnostic features that allow one to conclude whether they formed by biological or physical processes. However, Bradley found some grains shaped like whiskers or plates. These are only 50 to 150 nanometers long and have length to width ratios of 5 to 10 (so they are much longer than they are wide). According to Bradley and coworkers, these shapes are different from magnetites made by bacteria on Earth, which have length to width ratios of 3 or 4. Electron microscope image below shows two prominent (dark) magnetite whiskers in ALH 84001. The grains are much longer than they are wide. They are similar in size and shape to the suspected nanofossils. One magnetite grain is associated with a cavity in the carbonate (the white area).

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The whiskers have internal crystal structures that seem to differ from those of biologically-produced magnetite. Some have structures that resemble magnetites grown by a mechanism involving condensation of vapor to liquids, then to solid crystals. Others have distinctive lines along their lengths, as shown in the image on the right. This distinctive pattern is called a screw dislocation, and formed when the crystal grew. It is as if the crystal grew like a spiral staircase around a central post.

Magnetites formed by biological mechanisms have a high degree of crystal perfection. They do not contain screw dislocations or other structural defects. So, Bradley and his coworkers conclude that the magnetites in ALH 84001 were not made by martian microorganisms. Instead, the magnetites look like those formed at volcanic fumaroles, where they are deposited from a vapor.

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Example of a fumarole: Oblique aerial view taken in 1980 of an active 300-foot-diameter fumarole near Spirit Lake, Washington.

This requires a high temperature, 500 to 800 oC, much higher than life could have survived. Harvey and McSween had already concluded, on the basis of chemical compositions, that the carbonates had formed at such high temperatures, though that interpretation is hotly contested.

Most of the magnetites in ALH84001 do not contain the distinctive structures, so perhaps some formed by vapor deposition at some time in the rock's history, but most did not. The relation of the magnetite whiskers to suspected fossils is unknown. In other words, there may be more than one source of magnetite. Furthermore, there has not been enough study of biologically-produced magnetite to rule out the presence of whiskers and screw dislocations in them. A comprehensive study of magnetite in terrestrial microorganisms is needed before the magnetite case is closed.

As more data lead to additional interpretations, PSR Discoveries is committed to covering the continuing debate about evidence for life in a martian meteorite. See "Not quite a Meeting of the Minds" for an up-to-date summary of the scientific papers dealing with ALH 84001presented at the 28th Lunar and Planetary Science Conference held in Houston, Texas on March 17-21, 1996.

Becker, L., D. P. Glavin, and J. L. Bada, 1997, Polycyclic aromatic hydrocarbons (PAHs) in Antarctic Martian meteorites, carbonaceous chondrites, and polar Ice, Geochimica et Cosmochimica Acta., v. 61, no. 3 pp. Bradley J. P., R. P. Harvey, and H. Y. McSween Jr., 1996, Magnetite whiskers and platelets in the ALH84001 Martian meteorite: Evidence of vapor phase growth, Geochimica et Cosmochimica Acta., v. 60, no. 24, pp. 5149-5155. McDonald, G. D. and J. L. Bada, 1995, A search for endogenous amino acids in the Martian meteorite EETA79001, Geochimica et Cosmochimica Acta., v. 59, pp. 1179-1184.

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Wright, I. P., M. M. Grady, and C. T. Pillinger, 1997, An Investigation into the Association of Organic Compounds with Carbonates in ALH84001(abstract), Lunar and Planetary Science Conference XXVIII, p. 1589-1590.

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PSR Discoveries: Hot Idea: Evidence for lo...rature formation of carbonates in ALH84001

http://www.psrd.hawaii.edu/May97/LowTempCarb.html

posted May 22, 1997

Low-temperature Origin of Carbonates Consistent with Life in ALH84001 Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology, SOEST, University of Hawai'i

Much of the evidence for fossil life in martian meteorite ALH84001 is contained in globules, plates, and veins of chemically-complex carbonate minerals. For life to have existed in this rock, the carbonates must have formed at a temperature low enough to ensure survival, probably less than about 150 oC. Two recent papers in the same issue of the journal Science suggest that the carbonates formed at sufficiently low temperatures to permit life. The interpretations stem from measurements of the magnetic properties of minerals in ALH84001 (Kirschvink and others, 1997) and from the chemical and isotopic compositions of carbonates (Valley and others, 1997). Science is rarely so clear cut, of course, and not all data agree with a low-temperature origin. For example, see the PSR Discoveries article by Edward Scott about shock effects in ALH84001. References: Kirschvink, J. L., A. T. Maine, and H. Vali, 1997, Paleomagnetic evidence of a low-temperature origin of carbonate in the martian meteorite ALH84001, Science, vol. 276, p. 1629-1633. Valley, J. W., J. M. Eiler, C. M. Graham, E. K. Gibson, C. S. Romanek, and E.M. Stolper, 1997, Low-temperature carbonate concretions in the martian meteorite ALH84001: Evidence from stable isotopes and mineralogy, Science, vol 275, p. 1633-1638.

Rotated Magnets

The Earth has a substantial magnetic field, a feature that protects us from radiation by substantial solar flares. In fact, the Earth's magnetic field behaves as if there were a huge bar magnet inside the planet. The field arises from the motions inside the metallic, liquid outer core of the Earth; the motions are driven by the Earth's rotation and heat flowing from the inner, solid core. Other planets also have magnetic fields, but at present the fields of Mercury, the Moon, and Mars are quite weak, less than 1% of Earth's field. However, their magnetic fields might have been stronger in the past, and knowing how strong would be helpful in understanding the composition of their cores, how long the cores were molten, and the origin of planetary magnetic fields. But how can we measure a past magnetic field? The answer is simple: many rocks record the strength and direction of the magnetic field at the time of their formation. As hikers know, some minerals line up with the Earth's magnetic field, allowing them to determine which way they are walking. The original magnets were flakes of the mineral magnetite, Fe3O4, which is naturally magnetic. At high temperatures, above 580 oC, grains of magnetite are not magnetic. Thus,

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when, for example, a lava flow or other igneous rock solidfies, magnetite forms and is held in place by the other minerals in the solid rock, but is not magnetized until the temperature has fallen to 580 oC. At that temperature, called the Curie temperature, the magnetite becomes magnetized in the direction of the field at that time and place. The rock has recorded a lot of information about the magnetic field, and as long as the rock is not heated above the Curie point or is not chemically rotted by water and air, the record is preserved. Besides iron oxide, some iron sulfides also record the magnetic direction.

Mars might have had a magnetic field similar to Earth's, perhaps formed by fluids moving inside a molten core. If so, it would be as if a gigantic bar magnet sat inside the planet. In the diagram, the arrows represent the direction of the local magnetic field on the surface of Mars. (PSR Discoveries graphic.)

Magnetic directions in ALH84001

Joseph Kirschvink and his colleagues at the California Institute of Technology decided to measure the magnetic properties of a very small sample taken from an area of ALH84001 where the pyroxene crystals had been broken up. Their sample had a mass of only about 20 milligrams, yet they were able to separate grains and measure the strength and direction of the magnetism. The pyroxene fragments contained small inclusions of iron sulfide (FeS), which is probably the main carrier of the magnetism. Magnetic measurements are done in a special laboratory that is shielded from the Earth's strong field by about six tons of steel plates buried in the walls of the laboratory, and some high-tech devices closer to the magnetometer. It is like protecting samples from contamination during chemical analysis.

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Although the investigators have measured only a few grains so far, the results are startling, and shown schematically on the left. They found that the grains differed greatly in the direction of magnetization. So instead of all pointing in the same direction, they pointed in a wide range of directions, as indicated by the arrows. When the rock first crystallized, all parts would have been magnetically aligned. When the rock was damaged by an impact, some of the fragments rotated, causing separate grains to appear to line up in different directions. (PSR Discoveries graphic.)

The important point for understanding the temperature of formation of the carbonates is that they fill up the spaces in between the pyroxene fragments, so were deposited after the fragmentation. Most important, if the carbonates had come in hot, above the Curie temperature for iron sulfide, then the original magnetization would have been erased, and the magnetic minerals would have recorded the field at the time they cooled, and they all would have been aligned in the same direction. Since they are aligned every which way, the carbonate must have been no hotter than the Curie temperature, about 325 oC for FeS. This is still too hot for life to have existed, but it is an upper limit to how hot the rock could have been heated--the temperature was most likely much lower. Detailed analysis of their magnetic data lead Kirschvink and his colleagues to conclude that the magnetic grains were not heated above 110 oC, within the range in which life can exist. They suspect that further measurements may suggest a temperature as low as 40 oC.

The strength of the martian magnetic field

The Caltech magnetic measurements demonstrate the importance of ALH84001 beyond the issue of life on Mars. The strength of the field Kirschvink and coworkers measured is surprisingly high, suggesting that the magnetic field on Mars 4 to 4.5 billion years ago was approximately as strong as the magnetic field of the Earth now. This means that it is likely that Mars had a liquid, metallic (or at least electrically conducting) core very early in its history. And, because the younger SNC meteorites have only weak magnetization, the magnetic field decreased with time, perhaps because the core crystallized. We need many more measurements before a detailed picture of the evolution of the martian core and magnetic field are known with certainty, but this first glimpse is tantalizing.

Minerals and isotopes suggest low temperatures

The minerals present in a rock, the compositions and compositional homogeneity (uniformity) of those minerals, and the relative abundances of the oxygen isotopes in them are like little thermometers planted inside the rock. John Valley, with collaborators at his institution (University of Wisconsin, Madison), Caltech, the University of Edinburgh in Scotland, and the University of Georgia, tried to read the little thermometers in ALH84001.

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When viewed in an electron microscope, it is obvious that the carbonate globules are complicated. This photograph is a colorized image of the intensity of electrons bounced back from a polished surface of a sample of ALH 84001. The colors represent different minerals. Green is orthopyroxene (the silicate with iron and magnesium), blue is glassy plagioclase feldspar, and the various shades of red and orange are carbonate minerals with a range in chemical composition. (Photo courtesy of Ralph Harvey, Case Western Reserve University.)

Valley and colleagues paid special attention to the amazing complexity of the carbonates in the meteorite. To use those little thermometers inside rocks, the minerals must be in chemical equilibrium. This means that the minerals have reached an agreement with one another, an agreement impelled by the laws of thermodynamics, about how to distribute the elements composing them. If the minerals can be shown to be in equilibrium, then a vast storehouse of experimental data can be tapped to estimate the temperature of formation. Unfortunately, the carbonate minerals are not in equilibrium, making estimates of the temperature very difficult. As Valley and co-workers point out, carbonate minerals formed at low temperatures in shallow seas on Earth do not appear to be in equilibrium, and if you tried to calculate their temperature of formation you would conclude that the oceans are at a temperature of 500 oC! Valley and colleagues analyzed the compositions of the carbonates, and agree with Ralph Harvey (Case Western Reserve University) and Harry Y. McSween (University of Tennessee) about the compositional diversity of the carbonate globules and plates. However, they argue strongly that Harvey and McSween's conclusion that the carbonates formed at high tempertures (above 650 oC) is incorrect. Besides the problem of lack of equilibrium, Valley points out that if the temperature were high for even relatively short times, only days or even minutes, the minerals would have homogenized because elements are fairly mobile at high temperatures. Even more damaging, high temperatures would have caused assorted chemical reactions to take place, producing new minerals that are not observed in ALH84001.

Low-temperature Features

The chemical inhomogeneity of the carbonates is more consistent with a low-temperature origin. Elements move slower at low temperatures than at high temperatures, so minerals do not homogenize easily. Estimating the temperature precisely is still difficult, because of the lack of equilibrium. Nevertheless, a good guess can be made from the abundances of oxygen isotopes, specifically oxygen-16 and oxygen-18. The ratio of these isotopes is dependent on the formation temperature of different minerals. Valley and his associates measured the abundances of oxygen isotopes in ALH84001 carbonates using an ion microprobe, a high-tech device that can measure tiny quantities in tiny spots (about 30 micrometers across) on polished slabs of a rock. This gets around the extreme difficulty of physically separating small grains and analyzing by other techniques. It raises other analytical difficulties, but these seem to be manageable.

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The data obtained by Valley and co-workers are shown in the graph above. The oxygen isotope data are expressed as the amount of deviation of the ratio of O-18 to O-16 compared to a standard, almost always mean Earth ocean water. The important point is that the pyroxene has a low value of O-18/O-16, but the carbonates are higher, and much more scattered. If the pyroxene were in equilibrium with carbonate at a high temperature, as Harvey and McSween have argued, then the carbonate points should be in the range given by the purple horizontal bar, around 7 or 8 parts per thousands. Instead, the carbonate points lie far to the right of this range, consistent with formation at low temperature. (The ratio of oxygen isotopes is never exactly the same in co-existing minerals in a rock, even if they are in chemical equilibrium, mostly because the lighter isotope, oxygen-16, moves more readily than the heavier one. At high temperatures, the difference in their masses is less important than at lower temperatures, so the ratio is closer in co-existing minerals. However, at low temperature, the difference in the rates at which oxygen-18 and oxygen-16 move becomes more pronounced, leading to larger differences between minerals.) Although they cannot estimate the temperature precisely because of the lack of equilibrium, the similarity to carbonates formed at low temperature (certainly less than 300 oC ) on Earth leads Valley and his associates to conclude that the carbonates in martian meteorite ALH84001 formed at relatively low temperatures, too. Clearly, more work is needed to pin down the temperature more precisely-300 oC is too high for life to have existed.

Editor's note: The temperature at which the carbonates in ALH84001 formed is one of the most hotly debated issues about the evidence for fossils in the meteorite. For an opposing view, see PSR Discoveries article Shocked Carbonates may Spell N-o L-i-f-e in Martian Meteorite ALH84001 posted on May 22, 1997.

Kirschvink, J. L., A. T. Maine, and H. Vali, 1997, Paleomagnetic evidence of a low-temperature origin of carbonate in the martian meteorite ALH84001, Science, vol. 276, p. 1629-1633. Valley, J. W., J. M. Eiler, C. M. Graham, E. K. Gibson, C. S. Romanek, and E.M. Stolper, 1997, Low-temperature carbonate concretions in the martian meteorite ALH84001: Evidence from stable 5 of 6

isotopes and mineralogy, Science, vol 275, p. 1633-1638.

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PSR Discoveries: Lunar & Planetary Science Conference

http://www.psrd.hawaii.edu/Mar97/LPSCreport.html

posted March 31, 1997

Not Quite a Meeting of the Minds -- Report from the Lunar and Planetary Science Conference Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology

The annual Lunar and Planetary Science Conference was held at the Johnson Space Center in Houston from March 17 to March 21, 1997. The LPSC (as we all call it) is the largest conference devoted exclusively to planetary science, and certainly the most diverse in its coverage. The screens in one lecture room might be showing pictures of the amazingly intricate surface of Ganymede. Another might be displaying photographs of the result of an experiment designed to reproduce conditions in the cloud of dust and gas from which the planets formed. Still another could be showing images of lava flows on Venus. This was the 28th LPSC. For the first eight years of its existence, the conference was called the "Lunar Science Conference" and focused on results of studies of the data and lunar samples returned by the Apollo program. As the years went by, however, other planetary studies began to be reported at the conference, so in 1979 conference organizers added "planetary" to the name. They might need to add "biology" next as many researchers focus on martian meteorite ALH 84001 and its purported fossils. This was the year nanobacteria and biomarkers came to the LPSC.

In a break with our policy of reporting only on published research, PSR Discoveries has prepared this summary of the status of the scientific debate about life in martian meteorite ALH 84001. (See related October, 1996 article, "Life on Mars".) Rather than discuss the results of each of the talks given at the conference, I highlight the problem areas that need to be addressed before any consensus develops about the evidence for fossils in ALH 84001. If you'd like to see the range of topics discussed, take a look at the conference program. A wildly complicated rock

As more people look closely at ALH 84001, it is obvious that it is incredibly complicated. This is especially true of the carbonate occurrences in the rock. Scientists are focusing attention on the carbonates because that's where most of the evidence for life is found. They are complexly zoned in chemical and mineral compositions, and show intricate relationships among the all the minerals associated with the carbonates. There are somewhat large magnetite and sulfide mineral grains, and very tiny ones, and magnetite may occur in more than one location in the carbonate globules. The carbonate minerals occur in a variety of ways, including as round globules, flattened pancakes, vein fillings, and between fragmented pyroxene crystals (pyroxene is the main mineral in the rock).

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PSR Discoveries: Lunar & Planetary Science Conference

http://www.psrd.hawaii.edu/Mar97/LPSCreport.html

With all this complexity, it is no wonder no one agrees! It seemed clear to me when listening to the discussions at the conference that some of the disagreement stems from different investigators looking at different portions of the rock. It is as if they are climbing all over a huge elephant, each seeing a different part. One scientist sees a large eye and claims that the sample can see, but concludes that it cannot eat or smell as he found no mouth or nose. Another examines the tip of the trunk and suspects that the creature can smell, but appears to have no eyes. Another examines the surface of the trunk, concluding that the creature has a flexible leg that rarely touches the ground. None of them agrees with the other, and none of them has the entire picture. Another interesting problem arises because of differences in analytical approach. Some scientists emphasize the chemical compositions of the minerals, others the abundances of different isotopes of carbon, oxygen, or sulfur, while still others study the shapes of minerals and how they are intergrown. These different approaches tend to lead to different viewpoints. Only when all the observations are tied together will we begin to understand this wildly complicated rock from Mars.

Did the carbonates form hot or cold?

A very important issue is the temperature of formation of the carbonate minerals. Because the suspected fossils are associated with the carbonates, the carbonates must have formed at a suitably low temperature (less than 150 C) for life to have existed. If it could be proven that the carbonates formed at a high temperature, say above 250 C (some folks have suggested the carbonates formed at 500 to 700 C), then life could not have existed at the time of carbonate formation. Of course, proving a low-temperature origin does not prove that life existed in the rock, only that conditions were suitable for life. Geologists have devised several ways to determine the temperature at which a rock formed. In fact, we can also determine the pressure. However, most of the techniques we use to do this require two or more minerals to be in chemical equilibrium, which means they reached an agreement with each other about how much of each element or isotope they would contain. This level of agreement can be measured in the laboratory. Unfortunately, it appears that the minerals in ALH 84001 did not have time to reach an amicable agreement--they are not in equilibrium. This makes it very difficult to estimate the temperature of formation, and has led to the impressive disagreement in the estimated temperatures. The general compositions and isotopic abundances in the carbonates are like those in many carbonates on Earth, however, suggesting they formed at a relatively low temperature. On the other hand, few experiments have been done on rapidly crystallized carbonates at high temperature. The solution will come from more complicated ways of determining temperatures. These involve understanding the rates at which minerals form and the speed with which elements can move through minerals. This general type of problem is called kinetics, and is very important in all geologic studies. It is also an area where we lack sufficient data. So, for now, the temperature of formation of the carbonates is not known.

Contamination and related problems

The original paper by Dave McKay and his associates had carefully considered the problem of contamination of the sample with organic material while it was sitting on the ice in Antarctica. Nevertheless, some investigators, such as Luann Becker and Jeff Bader, have showed that contamination may be a serious problem (see their arguments summarized in Life on Mars--The Debate Continues). The meteorites found in Antarctica are preserved once they are buried in the ice, but when they emerge again at places like the Allan

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Hills, they may sit in puddles of water off and on for many years, allowing lots of time for organic chemicals and perhaps even bacteria to enter the rock. The problem may be quite serious, though more tests need to be done to determine the extent of contamination. The central problem is that the Earth is teeming with life and it seems to find places to live practically everywhere, even underground (see Life Underground, Dec. 96 article). Besides the opportunity for contamination in Antarctica, there is a chance for contamination in the curatorial facility at the Johnson Space Center, in spite of the careful procedures followed. In fact, since the discovery of possible life in ALH 84001, the curatorial laboratory has tightened its procedures to ensure that no contamination by organic chemicals will take place. A form of contamination may take place when samples are being prepared for analysis in a scanning electron microscope. The samples need to be coated with an electrical conductor to carry off the electrical charge added by the beam of electrons that is sprayed onto the sample. To coat the sample, gold, copper, or other metals are deposited onto the sample from a vapor. This process can produce unusual features on the rock surface, commonly called "sample preparation artifacts." The concern is that some of these artifacts might resemble the minute fossil-like objects in ALH 84001. To get around this, electron microscopes that do not require coatings are available, but have not yet been used extensively to study ALH 84001. Another way to get around the problem is to search for features similar to the ALH 84001 microfossils in other meteorites that do not contain fossils and that fell recently.

Do nanofossils exist on Earth?

The original paper by Dave McKay and coworkers compared the small fossil-like objects in ALH 84001 to very small microfossils found in some rocks on Earth. The problem is that these small features, called "nanobacteria" to emphasize that they are smaller than "microfossils," have not been proven to everyone's satisfaction to be fossils. The guidelines for determining that small features in ancient terrestial rocks are fossils have not been applied to the study of nanofossils as yet. (See Rules for Identifying Ancient Life, Oct. 96 article.) For example, no documented cases of cell division have been described. This lack of documentation means that the terrestrial nanobacteria used for comparison may not have been bacteria at all. On the other hand, they might be small bacteria, or the shrunken remains of bacteria, or parts of bacteria. Whatever the answer, detailed studies of both live bacteria that live in rocks and fossilized bacteria on Earth must be done before we can use them for comparison to the fossil-like features in ALH 84001.

When did the carbonates form?

One piece of evidence used to infer that ALH 84001 houses fossil life is the age of the carbonate globules in the rock. It was thought to be 3.6 billion years old. This is about the time when Mars was wettest, hence more likely to promote the development of life. However, another study concluded the age was only 1.3 billion years, much younger than the wettest period on Mars. This huge descrepancy will not be easy to understand because it is very difficult to date the carbonates. The minerals contain only small amounts of the elements used for dating by the usual methods, such as potassium-argon and rubidium-strontium, making the measurements extremely difficult on the small quantities available. In addition, some methods, such as rubidium-strontium, require that the minerals be separated from one another, a very difficult task in the complicated carbonate globules.

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At present, the only ages that seem to be well known is the original crystallization age, about 4.5 billion years, and an impact-heating age sometime between 3.8 and 4.05 billion years. If the carbonates were redistributed and heated by the impact event, then the carbonate age is the same as the impact age. But until the origin of the carbonate is worked out, this is not known.

Debate at the conference--during sessions, out in the hallways, over dinner--was vigorous, but far from vicious. Most of the proponents for one view or another (for example, high temperature versus low temperature origin for the carbonates), are quite sure they are right, which does arouse some passion! But all involved asked each other difficult questions that must be answered before we can say we know what happened to make ALH 84001 the way it is today. Some investigators will turn out to be right, others wrong, but there will be no losers. Everyone will have contributed to a better understanding of the geological history of Mars. Studies will help us understand carbonate formation, life in Antarctica, analytical techniques, nanofossils on Earth, climate on Mars, how to search for life on Mars, and how life evolved. And, the rock is interesting in its own right, independent of whether tiny organisms crawled around in it, and even independent of the origin of carbonates. ALH 84001 is a piece of the ancient crust of Mars, so it contains important information about the formation of the crust of that fascinating planet.

McKay, David S., and others, 1996, Search for Past Life on Mars: Possible Relic Biogenic Activity in Martian Meteorite ALH84001, Science, vol. 273, p. 924-930. Thousands of meteorites are collected in Antarctica through a program managed jointly by the National Science Foundation, Smithsonian Institution, and the National Aeronautics and Space Administration. Find out more from the Johnson Space Center and the Antarctic Search for Meteorites program.

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PSRD: Magma and Water on Mars

posted December 27, 2005

Magma and Water on Mars --- Martian meteorites tell us part of the fascinating story about when volcanoes erupted and water flowed.

Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology

Lars Borg (University of New Mexico) and Michael Drake (University of Arizona) synthesized available age data for Martian meteorites. Cosmochemists have determined when a variety of Martian igneous rocks crystallized and when their original minerals were altered by interaction with water. Igneous events occurred soon after the planet formed 4500 million years ago and continued to about 174 million years ago. Water affected the planet beginning within 100 million years of solar system formation and continued to less than 170 million years ago, perhaps to now. These observations tie in reasonably well with what we know from photogeologic studies of Mars, but we need more quantitative age determinations, implying sophisticated in situ age measurements and sample return missions. Reference: ●

Borg, Lars and Michael J. Drake (2005) A review of meteorite evidence for the timing of magmatism and of surface or near-surface liquid water on Mars. Journal of Geophysical Research, v. 110, doi: 10.1029/2005JE002402.

What Happened When?

Planetary scientists focus their attention on what geologic events shaped a planet, the relative sequence of events, and the absolute ages of those events. What happened? When did it happen? In the case of Mars, there are five major processes that manufactured its crust. One is the initial melting that formed a metallic core, primitive mantle, and initial crust above the mantle. This fundamental event involved assembling the planet from smaller planetesimals, widespread melting of the planet, and differentiation into its major zones. The second major process involves subsequent melting to produce a series of igneous rocks over billions of years. Some of the magmas erupted onto the surface; others were trapped in the crust to form what geologists call "intrusive" or "plutonic" rocks. The third major process involves water that carved channels and drainage networks, chemically reacted with the igneous rocks making up the crust, and transported sediments for deposition in low areas. The fourth process is impact. This was much more enthusiastic early in the history of the Solar System. In fact, the inner planets were assembled by planetesimals whacking into each other. Most of the craters in the ancient highlands of Mars were formed by high-velocity impacts of objects onto Mars. The final process is erosion and redistribution by wind, which may be the most important process now operating on the Red Planet.

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PSRD: Magma and Water on Mars

Processes that have shaped the Martian surface: initial planetary differentiation, igneous, aqueous, impact, wind.

Lars Borg and Mike Drake focus on the first three processes, with emphasis on the igneous and aqueous history of Mars. Their synthesis of age data allows us to divide Mars into broad age brackets on the basis of meteorite data.

Planet Construction Events

Cosmochemists can determine when a metallic core formed on Mars by measuring the isotopic abundances of hafnium (Hf) and tungsten (W). The short-lived isotope 182Hf decays into 182W. Because 182Hf has a half life of only 9 million years, it has very good time resolution for the early solar system. Furthermore, if metallic iron is present when a planet melts, tungsten concentrates in it, leaving behind a molten silicate mantle with very little tungsten, but the original amount of hafnium. If core formation took place while 182Hf was still around, 182W would build up in the rocky portion, allowing a measurement of when the metal separated from the silicate--the time of core formation. For more gory details, see PSRD article Hafnium, Tungsten, and the Differentiation of the Moon and Mars. The most recent measurement of the time of core formation was reported by Nicole Foley and colleagues (including Lars Borg). From hafnium and tungsten isotopes in Martian meteorites, Foley calculates that the Martian core formed 4556 (± 1) million years ago. That is no more than 12 million years after formation of the oldest objects in the Solar System, the calcium-aluminum rich inclusions in chondrites. http://www.psrd.hawaii.edu/Dec05/Magma-WaterOnMars.html (2 of 13)

PSRD: Magma and Water on Mars

After the metallic core formed, it may have taken a while for the surrounding molten or partly molten silicate (the primitive mantle) to solidify. How long that took is recorded, ironically, in the youngest Martian meteorites, the shergottites. Shergottites are basaltic rocks that formed as lava flows. In a paper published in 2003, Borg and his colleagues showed that the regions in the mantle that gave rise to the shergottites formed 4513 (± about 30) million years ago. Borg realized he could use both the long-lived and short-lived samarium (Sm) isotopes to define the age of the shergottite mantle sources by using three parameters to assess the age of differentiation of shergottite mantle sources. One is the initial ratio of neodymium-143 to neodymium-144 (143Nd/144Nd) in each shergottite, expressed as ε-143Nd, a measure of how the ratio deviates from the ratio in chondritic meteorites, in parts per ten thousand. 147Sm decays to 143Nd with a long half-life of 106 billion years. The other is the initial ratio of 142Nd/143Nd, or ε-142Nd, the deviation of the ratio from chondritic meteorites, again in parts per ten thousand. 142Nd is produced by the decay of 146Sm, which has a half life of only 103 million years. He also needed to know the ratio of 147Sm/144Nd in the source regions. Borg knew he could calculate lines of equal age (isochrons) for different times if he assumed three things: (1) the epsilon (ε) values were initially like those in chondrites, (2) formation of the sources involved formation of reservoirs with different Sm/Nd ratios, and (3) the sources melted only once more (when each shergottite magma formed). By plotting data from shergottites on the resulting complicated graph, Borg hoped to find the age of mantle differentiation. Adding new results obtained by Nicole Foley produces a slightly revised age of 4526 ± 20 million years.

Assuming that the parts of the Martian mantle formed early and remained unmelted until formation of the magmas that gave rise to the shergottites, Lars Borg showed that the shergottites lie on a single line that indicates formation of at least portions of the Martian mantle 4526 (± 20) million years ago.

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PSRD: Magma and Water on Mars

Taken at face value, it appears that the Martian core formed 30 million years before the overlying silicate mantle solidified (4556 My for core formation, 4526 for formation of the shergottite source regions in the mantle). However, Maud Boyet and Richard Carlson (Carnegie Institution of Washington) have made improved measurements of the ratio of 142Nd to 144Nd in chondrites, the basis for Sm-Nd age dating. They found that this ratio in chondrites is 20 parts per million lower than assumed previously. While not sounding like a big number, it translates to a big change in the estimate of the time of mantle formation on Mars. No detailed calculations have been reported, but it is possible that the time of mantle formation coincided with the time of core formation. Boyet and Carlson's result is new and its full ramifications are being explored by cosmochemists.

Early Crustal Events

One of the most famous Martian meteorites is also the oldest, ALH 84001. This meteorite is notorious for the proposed presence of tiny fossils and other biomarkers (see PSRD article Life on Mars?). While the debate about fossils in the rock has died down (with most investigators concluding that there are other explanations besides biology for the array of features in the rock), cosmochemists are paying more attention to how the original igneous rock formed. Larry Nyquist (Johnson Space Center) and coworkers found that the rock crystallized from a magma 4500 (± 130) million years ago. It must have formed in a body of magma intruded into the crust because it is composed mainly of one mineral, orthopyroxene. Such one-mineral rocks are called "cumulate" because they form by accumulation of minerals in a buried magma body.

LEFT: Large crystals of orthopyroxene in ALH 84001 show that this rock formed in an underground magma chamber on Mars. Dark areas are chromite (an oxide of chromium and iron). RIGHT: ALH 84001 formed about 4.5 billion years ago in a relatively large magma body inside the crust of Mars. Its high abundance of one mineral (orthopyroxene) indicates that this mineral must have accumulated in the magma, probably near the bottom of the magma body, eventually forming the original igneous rock with large crystals of orthopyroxene.

ALH 84001 contains intricate deposits of carbonate minerals. They reside in a variety of settings and textures, from interstitial crack fillings to conspicuously zoned clusters, semi-circular in cross-section, which have gotten the name "rosettes." (See PSRD article Carbonates in ALH 84001: Part of the Story of Water on Mars.) These assemblages are quite complicated. Carbonate compositions vary widely and they contain small grains of magnetite and sulfides. It would seem to be impossibly challenging to determine their age, but Lars Borg and colleagues at the Johnson Space Center managed to painstakingly separate different mineral phases by differential dissolution and then measured the isotopic composition of rubidium and strontium in the rock. The resulting Rb-Sr age is 3900 ± 40 million years. They also measured the age using lead isotopes, which gave the same value within experimental uncertainty, 4040 ± 100 million years. Apparently the carbonates precipitated from water that had soaked the rock about 3900 million years ago, showing that water could flow through the rock for at least a short time. Curiously, the rock has been bone dry since.

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PSRD: Magma and Water on Mars

Carbonate assemblages in ALH 84001. Carbonates are orange, clear, and dark. Surrounding material is low-calcium pyroxene (orthopyroxene). Field of view is about 0.4 mm across.

Borg and Drake summarize intriguing research done by Don Musselwhite and Mike Drake on understanding the ratio of xenon129 to xenon-132 (129Xe/132Xe) in the Martian atmosphere compared to Martian basalts. 129Xe is the decay product of another of those short-lived isotopes, in this case iodine-129. 129I has a half life of 16 million years. The Martian atmosphere has a high ratio of 129Xe/132Xe (2.4) compared to Martian meteorites (1.0 to 1.5). This indicates that early in Martian history iodine was separated from xenon. This had to happen before the radioactive 129I had decayed away, about 100 million years after synthesis of 129I in an exploding star before the Solar System formed. Borg and Drake summarize mechanisms for how iodine and xenon can be separated, concluding that the most promising involves water. Iodine is much more soluble in water than is xenon; in fact, it is 100 billion times more soluble. After initial differentiation is complete, iodine and xenon are driven from the interior by volcanism. If widespread bodies of water were present (oceans or big lakes), the iodine would dissolve in the water, leaving xenon in the atmosphere. Large impacts may have driven off this initial atmosphere, xenon included. The iodine remains in solution in the water. The water reacts with crustal rocks and drains inside the upper layers of the crust, and the 129I in it continues to decay, forming 129Xe and giving the rocks a high ratio of 129Xe to 132Xe. With time, the Xe is released into the atmosphere, giving a high 129Xe/132Xe ratio. If correct, this story implies that there was water present on the Martian surface soon after formation of the core and mantle.

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PSRD: Magma and Water on Mars

If oceans and other large bodies of water existed very early in Martian history, iodine released from erupting magma might have dissolved in them, thereby separating iodine from xenon. Painting courtesy of NASA.

The Billion Year-Plus Club

One group of seven Martian meteorites, the nakhlites, has an age of 1327 ± 39 million years. These are unusual lava flows in which pyroxene has accumulated. They also contain olivine, plagioclase feldspar, glass, calcium phosphates, and magnetite as igneous minerals. They might not all have formed at exactly the same time because their ages range from 1260 to 1370 million years, but they clearly represent a region of Mars that was volcanically active 1300 million years ago.

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PSRD: Magma and Water on Mars

Photomicrograph of a thin section of the Nakhla meteorite, viewed in polarized light. Blue and reddish crystals are olivine. All the others are high-calcium pyroxene (augite).

An interesting problem is the lack of meteorites between 1300 and 4500 million years old, in spite of much of the crust of Mars being constructed during that time. This may simply reflect how the meteorite delivery process (large impacts) samples the Martian surface and when it does so. For example, if an impact into a region of Mars dominated by rocks 3000 million years old took place 500 million years ago, or even 100 million years ago, the meteorites sent flying from Mars might have been swept up by the Earth and other planets long ago. Once on Earth, they would weather away or be buried, unavailable to us now.

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PSRD: Magma and Water on Mars

Hadriaca Patera is on the northeast rim of Hellas basin; north toward top of the image. Radial lava flows from this low, broad volcano can be seen between the patera and the channel of Dao Vallis. These are old areas from which we do not have meteorites.

The nakhlites also contain minerals formed by reaction with water. These include carbonates, sulfates, clay minerals, and iron oxides. One common form is called iddingsite, which contains clays, and assorted hydrated and not hydrated iron oxides (magnetite, maghemite, and ferrihydrite), all of which are common products of rock weathering. The weathering products typically form veins inside the meteorites, not unlike what happens during the initial stages of weathering of rocks on Earth, but usually make up about 1% of the rocks.

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PSRD: Magma and Water on Mars

Julie Stopar (University of Hawaii) created this mineral map of a thin section of the MIL 03346 nakhlite. Colors correspond to different minerals. Most are igneous, but veins of altered rock (blue in the map) and patches of sulfate minerals (red) occur. The alteration zones are like those dated in other nakhlites.

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PSRD: Magma and Water on Mars

Alteration vein (reddish brown) inside an olivine crystal in MIL 03346 nakhlite is composed of weathered olivine with sulfur (S) and chlorine (Cl) added.

Age dating experts at the Johnson Space Center and the University of Arizona attempted the very difficult task of dating the products of aqueous alteration. The Johnson Space Center group did differential dissolution experiments to get at least two points on a rubidium-strontium diagram. The resulting ages were 614 to 679 million years. The Arizona group tried to use potassium-argon dating. They found a wide range in ages, from about 100 to 670 million years. Those less than 600 million years might have lost argon because of heating, perhaps when the meteorites were blasted off Mars. The best guess for the age of alteration in the nakhlites is 633 ± 23 million years.

Rock Youngsters

The youngest Martian meteorites are basaltic rocks that formed in lava flows, the shergottites. They consist of pyroxene (two types, pigeonite and augite), plagioclase feldspar (transformed by impact into a glass), and assorted accessory minerals. They are the most abundant Martian rocks in our collections on Earth.

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PSRD: Magma and Water on Mars

Tara Hicks-Johnson (University of Hawaii) created this mineral map of a thin section of the Shergotty Martian meteorite. It consists mostly of pyroxene and plagioclase feldspar, the main minerals in planetary basalts.

The shergottites range in age from 165 to 575 million years old. However, there are clear groupings of ages: 575 My, 474 My, 332 My, and 175 My. Considering the poor sampling of Mars that meteorites give us, this indicates quite a bit of recent volcanism. Equally important, isotopic measurements determine the initial ratio of strontium-87 to strontium-86 (87Sr/86Sr) in the lava flows, which is a fingerprint of the chemical characteristics of the places in the mantle that melted to make the magmas. These studies indicate that the 10 shergottites in the 175 My group represent no fewer than 10 separate regions in the mantle, hence 10 separate melting events. If we take the full range of apparent ages, the maximum time over which these basalts erupted is 42 million years. That translates to one eruption every 4.2 million years. If the spread in ages is represented by the standard deviation of the average in the ages, 4 million years, then there is an eruption every 0.4 million years. Not too active by terrestrial standards, but considering sampling biases, it shows clearly that Mars was volcanically active within the past 175 million years, suggesting there could be some activity very recently. All the shergottites contain the products of aqueous alteration, including sulfates and carbonates. They are very low in abundance, and no direct age determinations have been made on them. However, they occur in rocks that are all younger than 575 million years, many 175 million years old, so the aqueous alteration must have occurred during the past 175 million years. Hence, there is recent volcanism on Mars and recent aqueous alteration. Water may not be flowing across the dusty surface as vigorously as it did billions of years ago, but there is enough of it around to react with basalts, at least in some places.

More Datable and Partly Rotted Rocks Await Us on Mars

The meteorite research Lars Borg and Mike Drake summarize shows the value of being able to look closely at rocks, identify igneous and alteration minerals in them, and to make painstaking isotopic measurements and age determinations. The only

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PSRD: Magma and Water on Mars

drawback is that the meteorites do not contain the complete story. There might be a bias in how rocks are blasted off Mars and delivered to us stuck here on Earth. Rocks that have been extensively altered by water might not be able to survive such a violent launch process. Meteorites from the large area with ages between 1300 and 4500 million years might not have been sent our way at the right time that we find them on Earth now. We need more samples.

Well, there are plenty of samples. It's just that they are on Mars, which requires that we get to them. Many studies will be possible with the next generation of rovers and analytical instruments on them. Certainly specialized instruments will be able to identify many of the weathering products. However, dating the rocks to determine when they formed and when they were altered by water requires another step in our analytical capabilities. Some cosmochemists are working on that problem, hoping to be able to date rocks with an accuracy of 10 or 20%. On the other hand, many cosmochemists are hoping that well-chosen samples are returned from Mars so they can be studied by the array of high-precision isotopic techniques available in terrestrial laboratories like Lars Borg's at the University of New Mexico.

LINKS OPEN IN A NEW WINDOW.

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Borg, L. and M. J. Drake (2005) A review of meteorite evidence for the timing of magmatism and of surface or nearsurface liquid water on Mars. Journal of Geophysical Research, v. 110, doi: 10.1029/2005JE002402. Borg, L. E., L. E. Nyquist, H. Wiesmann, and Y. Reese (2003) The age of Dar al Gani 476 and the differentiation history of the Martian meteorites inferred from their Rb-Sr, Sm-Nd, and Lu-Hf isotopic systematics. Geochim. Cosmochim. Acta, v. 67, p. 3519– 3536. Boyet, M. and R. W. Carlson (2005) New 146Sm-142Nd Constraints on Moon Formation and Early Silicate Planetary Differentiation. Meteoritics & Planetary Science, v. 40, p.5246. Corrigan, C. M.(2004) Carbonates in ALH84001: Part of the Story of Water on Mars. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/July04/carbonatesALH84001.html Foley, C. Nicole, M. Wadhwa, L. E. Borg, P. E. Janney, R. Hines, and T. L. Grove (2005) The early differentiation history of Mars from 182W and 142Nd isotope systematics in the SNC meteorites. Geochimica et Cosmochemica Acta, vol 69, p. 4557-4571. Martel, L.M.V. (2001) Outflow Channels May Make a Case for a Bygone Ocean on Mars. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/June01/MarsChryse.html Martel, L.M.V. (2003) Ancient Floodwaters and Seas on Mars. Planetary Science Research Discoveries. http://www. psrd.hawaii.edu/July03/MartianSea.html

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Taylor, G. J. (2000) Liquid Water on Mars: The Story from Meteorites. Planetary Science Research Discoveries. http:// www.psrd.hawaii.edu/May00/wetMars.html Taylor, G. J. (2003) Hafnium, Tungsten, and the Differentiation of the Moon and Mars. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/Nov03/Hf-W.html Taylor, G. J. (2004) Multifarious Martian Mantle. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/ June04/martianMantle.html Taylor, G. J. (2005) Recent Activity on Mars: Fire and Ice. Planetary Science Research Discoveries. http://www.psrd. hawaii.edu/Jan05/MarsRecently.html

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http://www.psrd.hawaii.edu/Dec05/Magma-WaterOnMars.html (13 of 13)

PSRD:A Primordial and Complicated Ocean of Magma on Mars

posted March 31, 2006

A Primordial and Complicated Ocean of Magma on Mars --- Geophysical and geochemical calculations indicate that total melting of Mars during its formation could have led to large-scale heterogeneities in its mantle.

Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology

It seems almost certain that the Moon was surrounded by an ocean of magma when it formed. This important

idea has been applied to the other terrestrial planets and even to asteroids. Linda (Lindy) Elkins-Tanton and

colleagues Mark Parmentier, Paul Hess, and Sarah Zaranek at Brown University, and Lars Borg and David

Draper (University of New Mexico) have examined the chemical and physical consequences of magma ocean

crystallization on Mars. Elkins-Tanton has focused on the fate of the pile of crystals created during

solidification of a magma ocean over a thousand kilometers thick. Crystallization causes the minerals that form

first to lie beneath those formed later. The deepest minerals are also less dense than the overlying minerals. This

is an unstable situation: the low-density rocks would have a tendency to rise while the high-density rocks would

have a tendency to sink. Although we think of rocks as solid and hard, when hot and under pressure, they flow

like liquids. They do not flow fast, but they do flow like ultra-gooey liquids (about a factor of 100 million

billion times gooier than ketchup at room temperature). Thus, the heavy layers sink and the light layers rise,

producing a complicated Martian mantle with chemical characteristics like those cosmochemists infer from

studies of Martian meteorites. The sinking of relatively cool rocks from the top of the crystallized pile cools the

boundary between the metallic core and the mantle, causing motions inside the core to produce the early, strong

magnetic field of Mars.

References:

z Borg, L. E. and D. S. Draper (2003) A petrogenetic model for the origin and compositional variation of the

martian basaltic meteorites. Meteoritics and Planetary Science, v. 38, p. 1713-1731.

z Elkins-Tanton, L. T., E. M. Parmentier, and P. C. Hess (2003) Magma ocean fractional crystallization and

cumulate overturn in terrestrial planets: Implications for Mars. Meteoritics and Planetary Science, v. 38, p.

1753-1771.

z Elkins-Tanton, L. T., E. M. Parmentier, and P. C. Hess (2005) Possible formation of ancient crust on Mars

through magma ocean processes. Journal of Geophysical Research, v. 110, E12S01,

doi:10.1029/2005JE002480.

z Elkins-Tanton, L. T., S. E. Zaranek, E. M. Parmentier, and P. C. Hess (2005) Early magnetic field and

magmatic activity on Mars from magma ocean cumulate overturn. Earth and Planetary Science Letters, v. 236,

p. 1-12.

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PSRD:A Primordial and Complicated Ocean of Magma on Mars

An Exceptionally Brief History of Mars

Mars was a geologically happenin? place during its first billion years of existence, and particularly during its first 50 million years or so. Substantial early melting caused formation of a metallic core, silicate mantle, and at least half of the rocky crust within 50 million years after the planet formed. Metallic iron dribbled to the center to make a metallic core within about 15 million years after the formation of the solar system. These early events are recorded by the isotopes in Martian meteorites [see PSRD article Magma and Water on Mars], by the presence of highly-cratered terrain in the highlands and underlying the smooth northern plains, and by magnetized regions in the most ancient terrain. This early activity may have been driven by rapid accretion of Mars from countless smaller objects, a manufacturing process that might have taken only a million years after the beginning of the solar system. After the initial pulse of melting, much of the igneous and tectonic activity focused around the Tharsis region, home of prominent volcanoes on Mars. Tharsis is a huge bulge in the crust, pushed up by forces below and decorated with volcanoes and lava flows. Magmas may have heated the crust, releasing water to form vast valley networks. Volcanism spewed large quantities of water and carbon dioxide into the atmosphere, further helping to erode the surface, though how much was eroded by rain and for what length of time water could have been stable is not known. Much of the construction of Tharsis was complete by the end of the Noachian period, well over 3 billion years ago. Volcanism and water-related processes have been intermittent since then. Because so much happened during the first billion years, cosmochemists are particularly interested in it. This applies especially to the first 50 million years, when the planet underwent its initial differentiation into metallic core, silicate mantle, and primary crust. Because the planet probably formed rapidly, it might have melted like the Moon did. This early melting would have set the stage for the subsequent geologic history of Mars. In other words, it?s a big deal, hence the interest in this event by Lars Borg and Dave Draper, and Lindy Elkins-Tanton and her colleagues.

The Moon as a Model

Lunar samples brought to Earth by Apollo astronauts changed the way we look at the early history of the Moon and inner planets. Before astronauts collected the first samples, we knew little about the Moon?s chemical composition. We did know from careful geologic mapping using telescopic and then Lunar Orbiter photography the relative sequence of geological events that shaped its surface, but we did not know the ages of rocks in the highlands or when lavas made the maria. In fact, we did not know for sure that the lunar maria (the dark areas on Moon) were composed of lava flows. Cosmochemists were shocked when they looked at the samples returned by the first piloted lunar landing, Apollo 11. In the charcoal gray grit scooped up by Neil Armstrong and Buzz Aldrin were white rock fragments a few millimeters across. They were clearly different from the other rock fragments, which were either pieces of lava flows demolished by impacts or mixtures of rocks (breccias) formed when impacts compacted lunar regolith. Everybody who examined the soil samples in detail noted the white fragments and showed that they were quite different from the basalts that make up the Sea of Tranquility. The white fragments were made almost entirely of one mineral, plagioclase feldspar. Such rocks are called "anorthosite." In a bold leap, John Wood, then at the Smithsonian Astrophysical Observatory, suggested that the anorthosite fragments were tossed to the Apollo 11 landing site by impacts in the highlands. He claimed that not only were the nearby highlands composed of anorthosite, but all the highlands were. How could such a large area be composed of one rock type? Wood took another bold leap and said that when the Moon formed it melted, producing a huge ocean of magma around it. Plagioclase, a low-density mineral, floated to the top, while denser minerals sank. The millimeter-sized anorthosite rock fragments were pieces of the primary lunar crust.

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PSRD:A Primordial and Complicated Ocean of Magma on Mars

The photograph on the left shows a collection of 2 to 4 millimeter rock fragments sieved from an Apollo 11 soil sample by John Wood and his colleagues at the Smithsonian Astrophysical Observatory in 1969. Most fragments are dark basalts or impact breccias, but pieces of white, feldspar-rich rock are also present (near the millimeter-scale bar). These sparked Wood?s imagination, leading to the idea of the lunar magma ocean.

In the ocean of magma covering the baby Moon, lightweight minerals floated and heavy ones sank. The lighter minerals formed the crust of the Moon.

The magma ocean was an imaginative idea, and there was some other evidence for it. The Apollo 11 basalts provided one bit of evidence. Cosmochemists found that some chemical characteristics of mare basalts were complementary to those of anorthosites. This suggested that the regions in the deep interior of the Moon where the basalt lavas formed by partial melting were part of the same magma from which the anorthosites formed. Since the mare basalts formed at depths of hundreds of kilometers, the magma must have been hundreds of kilometers thick. A prominent example of such complementary chemical features is the abundance of europium, which is depleted in mare basalts and enriched in anorthosite. Further evidence was provided by the presence of abundant anorthosite at the Apollo 16 landing site in the lunar highlands.

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PSRD:A Primordial and Complicated Ocean of Magma on Mars

Plot of the abundances of the rare earth elements for lunar anorthosites, mare basalts, and KREEP. Note the large, positive europium (Eu) anomaly in the anorthosite. The mare basalts have negative Eu anomalies, indicating that the minerals making up their source regions in the lunar mantle formed in a magma from which plagioclase feldspar (the main constituent in anorthosite) had been removed. The negative anomaly in KREEP is even more severe, consistent with it representing the last dregs of magma ocean crystallization.

Photograph of the first large anorthosite sample returned from the Moon, rock 15415, collected during the Apollo 15 mission. Apollo 16 returned lots of anorthosite samples.

In spite of this early success, we still lacked the most basic evidence: proof that the ancient highlands are made of anorthosite. The Clementine mission (1994) and my colleague Paul Lucey (University of Hawaii) solved this problem. Clementine snapped pictures of the entire lunar surface. Lucey, using Apollo landing sites and returned soils as ground truth, figured out how to convert the intensity of reflected light to the concentration of FeO [see PSRD article Moonbeams and Elements]. Anorthosite is composed chiefly of feldspar, which has very little FeO. It does contain some minerals with FeO, so anorthosite might contain a few percent FeO. The maps produced from the Clementine data show huge regions of the highlands, especially on the farside, that contain between 2 and 6 wt% FeO, with an average about 4 wt%, as predicted by the magma ocean hypothesis. Measurements by the Lunar Prospector mission confirmed the Clementine measurements.

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PSRD:A Primordial and Complicated Ocean of Magma on Mars

The maps, above, show the concentration of FeO on the Moon, as determined from Clementine data using a technique developed by Paul Lucey (University of Hawai'i) and updated by Jeff Gillis-Davis (previously at Washington University in St. Louis and now at the University of Hawai'i), Brad Jolliff, and Randy Korotev (both at Washington University in St. Louis). Low FeO translates into high Al2O3, consistent with the presence of lots of anorthosite, as predicted by the magma ocean hypothesis.

Once the magma ocean idea became entrenched in our thinking, cosmochemists began to examine the processes that could have operated in it. They began with knowledge gained from studies of layered intrusions on Earth. These are large bodies of magma that solidified far beneath the surface. As they cooled, minerals formed distinctive layers, some looking a bit like sediments deposited by rivers, but in this case the rivers were masses of swirling magma. However, a big difference was in scale. Terrestrial layered intrusions perhaps reach 8 kilometers in thickness and range in area from 100 km2 (Skaergård, E. Greendland for example) to 66,000 km2 (Bushveld, S. Africa). The lunar magma ocean was globe encircling and hundreds of kilometers deep. The terrestrial examples took us only so far. Cosmochemists depicted numerous processes operating in the magma ocean. It would have crystallized substantially at its base because of the higher pressure there. It would have been vigorously convecting, preventing minerals from settling until the magma became choked with crystals. Once plagioclase floated to the top, it would begin to form a crust, creating what John Longhi (now at Columbia University) called rockbergs. There would have been reactions at the margins of the rockbergs, changing the composition of the magma. As the base solidified, residual, highly-evolved magma would have oozed upwards, mixing with the overlying magma, or even traveled through the magma all the way to the growing crust to crystallize there. As crystallization proceeded the magma became richer in a group of elements that do not readily concentrate in the major minerals forming from the magma (olivine, pyroxene, plagioclase, and ilmenite). They became greatly enriched in the leftover magma ocean. These elements include potassium (K), rare earth elements (REE), and phosphorus (P). This leftover stuff became a prominent constituent of many rocks returned from the Moon, and received the nickname KREEP. Paul Warren and John Wasson (University of California, Los Angeles) called the very last leftover magma "urKREEP," using a German prefix that means "primary." UrKREEP was also rich in thorium (Th), which allowed the Lunar Prospector mission to use its gamma ray spectrometer to map its distribution. For obscure reasons, thorium is concentrated on the lunar nearside.

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PSRD:A Primordial and Complicated Ocean of Magma on Mars

This is John Longhi's 1979 sketch of the processes operating in the lunar magma ocean. Even this complicated diagram is probably an oversimplified view of what happens in a huge body of magma.

Thorium (Th) is not distributed uniformly on the Moon. Somehow it and other elements that are not readily incorporated into major minerals ended up concentrated in one region of the Moon.

Once the magma ocean had crystallized, it would have been denser at the bottom than at the top because the first minerals to form would have had higher magnesium to iron ratios. For example, olivine ranges in density from 3.2 if it is Mg2SiO4 to 4.3 if it is Fe2SiO4. The density of the olivine crystallizing from the magma ocean did not reach either extreme value, but there was a significant difference between the lower and upper portions of the accumulated pile of crystals. Pyroxene behaves the same way, and once ilmenite (FeTiO3) crystallized (density from 4.5 to 5) the rocks on top would have been quite dense, resulting in a tendency for the whole pile to overturn. Although mostly solid, the hot rocks would have been ductile enough to move slowly upwards and downwards, forming a more stable, but complicated lunar interior.

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PSRD:A Primordial and Complicated Ocean of Magma on Mars

Dense, titanium-rich rock might have sank through the underlying rock in the lunar mantle, while low-density rocks rose, scrambling the lunar mantle.

Cosmochemists have investigated the consequences of a magma ocean encompassing the primitive Earth and asteroids. Detailed models for Mars have only recently been discussed. This work is driven in part by our growing understanding of the chemical characteristics of Martian meteorites.

A Geochemical View of the Martian Magma Ocean

The Moon has a primary crust made of anorthosite, formed by plagioclase feldspar floating in its magma ocean. We do not observe this on Mars, but should we expect to see it? If not, what should we see? In other words, how do we test whether there was a magma ocean or not on Mars? Borg and Draper, and Elkins-Tanton and her colleagues have focused on the properties of the Martian mantle we infer from Martian meteorites. Although we have only about 30 of these important rocks and all but one are among the youngest rocks on Mars (all less than 1.3 billion years old), they still contain high-fidelity information about the mantle and when it formed. As detailed in the PSRD article The Multifarious Martian Mantle, the shergottite group of Martian meteorites indicates that there are at least two distinct reservoirs in the mantle. They have the characteristics outlined in the table below. Enriched Reservoir

high La/Yb low Sm/Nd (-εNd) high Rb (high 87Sr/86Sr) oxidized

Depleted Reservoir

low La/Yb high Sm/Nd (+εNd) low Rb (low 87Sr/86Sr) reduced

There are other reservoirs as well, such as those that manufactured the nakhlite group of Martian meteorites and another that produced the 4.5 billion year old ALH 84001 meteorite, but the shergottite-producing regions give cosmochemists modeling a Martian magma ocean something to use as a test. Another important feature is that the regions of the mantle in which the Martian meteorites formed are depleted in aluminum compared to terrestrial basalts. Magma ocean models must account for that as well.

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PSRD:A Primordial and Complicated Ocean of Magma on Mars Mars is much larger than the Moon, causing the pressure inside Mars to be higher than in the Moon. This greatly affects calculations of crystallization in a magma ocean on Mars compared to the Moon--the deeper the magma ocean, the higher the pressure in it. In a Martian magma ocean deeper than several hundred kilometers, high pressure minerals such as garnet form. This greatly changes the course of crystallization compared to the lunar magma ocean. Garnet can contain a hefty amount of aluminum, which ends up deposited at depth. It is not available to make an anorthosite crust as on the Moon. Furthermore, garnet and other minerals affect the concentrations of trace elements in the magma ocean as it crystallizes. Fortunately, all this can be modeled mathematically using experimental data as a basis. When modeling the lunar magma ocean, cosmochemists used a thorough knowledge of the order in which minerals crystallize in magma at low pressure. There are good computer programs available that enable these calculations. Unfortunately, such computer programs are not available for higher-pressures systems like those inside Mars. Besides that, experimental coverage at the right range of pressures is not thorough enough. This means that the crystallization sequence is not known as well as it is for the Moon. As a result, Borg and Draper had to make some reasonable guesses based on cosmochemical savvy. Fortunately, Dave Draper, Carl Agee (University of New Mexico), and their colleagues have been working to fill in gaps in the experimental database at high pressure. Many experimental studies done at high temperature and pressure had used a composition not close enough to what cosmochemists think Mars has. (Of course, that composition is a bit uncertain, too!) The major difference is the dissimilarity in the concentration of iron oxide compared to Earth and the Allende meteorite, the subjects of other experimental studies. Draper is a skilled experimentalist. He takes samples of silicate material (rocky stuff) and puts them in an experimental apparatus that heats and squeezes the sample. The pressure in his experiments reached 150,000 times that at the surface of the Earth at a temperature as high as 2000 Celsius. Dave Draper and his colleagues used this hydraulic device for experiments at high pressure and high temperature. The sample is contained in between ceramic octahedra located inside tungsten carbide cubes, which are placed inside a cylindrical module. The press then squeezes it all to the desired pressure. This photo was taken in 2002 when the apparatus was located at the Johnson Space Center. Standing next to the press is Jana Berlin, an undergraduate participant in the Summer Intern Program run by the Lunar and Planetary Institute [website]. More details of the apparatus, which is now at the University of New Mexico, can be found at the high-pressure laboratory's [website]. The lab was moved to UNM from the Johnson Space Center. (The construction of the new lab and moving the equipment makes an interesting story.)

The experiments filled in gaps in our knowledge of what minerals crystallize at different pressures inside Mars. The picture is not complete, but it is a reasonably detailed sketch. It allowed Borg and Draper to determine the crystallization sequence in a Martian magma ocean, hence to trace the way elemental concentrations changed as each layer crystallized from the magma. They studied three cases: shallow, medium, and deep magma oceans (see diagram below).

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PSRD:A Primordial and Complicated Ocean of Magma on Mars

Cartoon showing crystallization sequences, on the left, of the lunar magma ocean (calculated by Greg Snyder, then at the University of Tennessee) and, on the right, three Martian magma oceans with different depths, hence pressures. Higher pressure in deeper oceans results in different distributions of the elements. Borg and Draper do not intend these diagrams to represent the precise layering developed in a magma ocean because accumulated layers may not be uniform around the entire globe and because, as discussed below, they might rise or sink depending on their densities.

Borg and Draper calculated how the concentrations of major elements changed with magma ocean crystallization and accumulation of the crystallizing minerals. In a separate set of calculations they assessed how the concentrations of trace elements changed in the layers of accumulating minerals and in the leftover magma. The leftover magma, which cosmochemists typically dub "residual magma" or "trapped liquid," plays a critical role in the chemical characteristics of the magma produced by subsequent melting of cumulate layers. Borg and Draper compared these calculated magmas to the compositions of Martian meteorites or the estimated compositions of the primary magmas of the Martian meteorite lava flows. The Martian meteorites are characterized by a depletion of aluminum oxide (Al2O3) compared to typical basalts on Earth. This is usually expressed as the ratio CaO/Al2O3. The calculations indicate that Martian meteorite magmas cannot be made by cumulate rocks that formed in a shallow, low-pressure magma ocean. Such conditions lead to lower CaO/Al2O3 ratios than observed in the Martian meteorites. This suggests to Borg and http://www.psrd.hawaii.edu/Mar06/mars_magmaOcean.html

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PSRD:A Primordial and Complicated Ocean of Magma on Mars Draper that garnet (formed at high pressure) must have crystallized to remove aluminum from the magma. Recent results from the Mars Exploration Rovers suggest that not all Martian rocks have high CaO/Al2O3. These might represent magmas formed in portions of the mantle that were not depleted in aluminum. Alternatively, they might have formed much longer ago than were the Martian meteorite magmas. Their formation might have removed more aluminum than calcium from the mantle, leaving behind a depleted mantle containing high CaO/Al2O3 from which the young Martian meteorites formed.

This graph shows how the ratio of calcium oxide to aluminum oxide would vary as a function of the ratio of magnesium oxide to magnesium oxide plus iron oxide (Mg#) during partial melting of cumulates from a Martian magma ocean. The letters denote assorted estimates of the magmas that gave rise to some of the Martian meteorites. Cumulates formed at low pressure do not produce magmas (bottom line) with sufficiently high CaO/Al2O3, but those formed at high pressure do. This suggests that the magma ocean on Mars was deep, allowing formation of garnet, which sops up aluminum. This leaves the overlying rock depleted in aluminum and having high CaO/Al2O3.

The trace element calculations were done only for the deep magma ocean case because that is required to explain the CaO/Al2O3 ratio. Borg and Draper show that magmas representing the distinctive reservoirs discussed above could form by various combinations of cumulate minerals and trapped liquids. The results are not perfect, perhaps indicating additional complexities need to be taken into account. For example, potassium in calculated magmas does not match those of Martian meteorites. This might indicate that a potassium-bearing mineral (such as phlogopite) formed and affected trace element distributions. However, Borg and Draper tested this idea, too. They found no combination of minerals that could explain all the discrepancies at once. If potassium worked out, then tantalum, for example, did not. After publication of the 2003 paper by Borg and Draper, Dave Draper and his colleague Carl Agee reported on a series of high-pressure experiments using H-chondrites as starting materials. These experiments led them to conclude that the bulk composition of the rocky portion of Mars (hence of the original magma ocean) may have contained less iron oxide than cosmochemists have thought previously. This led Draper to run his magma ocean calculations again, using the slightly lower iron oxide concentration (12 or 14 wt%, rather than 18 wt%). The match with Martian meteorites is even better than before. It still explains the high ratio of CaO to Al2O3. And, the less iron-rich mantle reproduces a property called the moment of inertia factor. This factor was determined with great accuracy from measurements during the Pathfinder mission of how the pole of Mars changes the direction it is pointing. It assesses the distribution of mass inside the planet. The previous calculations with

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PSRD:A Primordial and Complicated Ocean of Magma on Mars higher iron oxide also reproduce the moment of inertia factor, but this new calculation is even closer. None of these calculations prove that a magma ocean existed on Mars, but they show that it is feasible. And complicated. The geochemical calculations are quite intricate, yet oversimplify what must have happened in a Martian magma ocean and afterwards. Some of that afterwards story is told by Lindy Elkins-Tanton and her colleagues.

A Geophysical View of the Martian Magma Ocean

Planets are dynamic. As a Martian magma ocean crystallized, the lowermost layers would have lower density than did layers higher up. This gradient in density is caused by the compositional changes in the magma as it evolves during crystallization. One major factor is that the ratio of iron to magnesium in olivine and pyroxene increases as the magma crystallizes. Fe/Mg decreases with depth in the crystallized magma ocean. There would be a tendency for lower layers to rise, the way blobs rise from the bottom of those sometimes fashionable but never classy lava lamps. Lindy Elkins-Tanton and her colleagues have studied the nature of this overturn. The first step in calculating the overturn of a magma ocean cumulate pile is to make the pile in the first place. This is much more complicated than assessing an order of crystallization. You have to worry about how the crystals separate and whether they are swept up in the convective flow that must accompany cooling of a huge, globe-encircling body of magma. There is an exquisite body of literature evaluating crystallization in magma bodies. It describes fluid dynamic theory, field observations of large layered igneous intrusions, and laboratory experiments. In spite of all this research, our understanding is not complete. Nevertheless, it allowed ElkinsTanton to calculate two reasonable cases for the distribution of minerals and the density variation with depth after magma ocean crystallization. These are summarized in the diagram below.

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PSRD:A Primordial and Complicated Ocean of Magma on Mars

Estimate of how the density of the Martian mantle varied with depth right after magma ocean crystallization and after it overturned. Elkins-Tanton tested two cases. In one case crystals settle in the general order in which they crystallize and with the densest at the bottom. In the other case, some of the garnet accumulates at the pressure at which it is dense enough to settle, but the other minerals (olivine and pyroxene) are buoyant and remain entrained in the remaining liquid. The garnet forms a dense layer in the middle of the pre-overturn mantle. The density variations (hence the abundances of minerals with depth) would probably not have been uniform around the entire planet.

Modeling the crystallization of a magma ocean is tricky. So is modeling the overturn of the products of that crystallization. The rate and duration of overturn is dependent on other factors besides density. A calculation http://www.psrd.hawaii.edu/Mar06/mars_magmaOcean.html

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PSRD:A Primordial and Complicated Ocean of Magma on Mars must include consideration of temperature because density, crystal structure, and compressibility all depend on temperature. The extent to which a mineral becomes denser with pressure (compressibility) is also important. So is the stability of high pressure minerals: as a portion of mantle rises, some minerals transform to lower density forms, impeding ascent. On top of all that, the viscosity of the flowing mantle rock changes with pressure and temperature. It is one complicated physics problem! Lindy Elkin-Tanton's colleague at Brown University, Sarah Zaranek, produced highly illustrative visualizations of the overturn (see below). They show that when the mantle is finished overturning, the result is not simply an orderly rearrangement of the layers. It is a messy process that produces lateral heterogeneities, unless the viscosity is constant throughout the mantle and throughout the overturn event (which was almost certainly not the case). Lateral variations in composition would produce complicated variations in rock compositions on Mars, which is observed by instruments on the Mars Odyssey spacecraft.

These colorful models of the Martian mantle show evolution of density stratification and movement (colored layers) after crystallization. As overturn occurs (see lefthand side), the dense upper layers sink and the less dense lower layers rise. The final product is a complicated mixture of materials. On the righthand side of the figure, we can track the initial location of a layer and see where it ends up in the complicated Martian mantle.

The two movies, linked below, show what happens with overturning of a cumulate Martian mantle with temperature- and pressure-dependent flow and deformation of the materials. These are animations of the bottom row of figures shown above.

+ View Movie 1 Initial layering in density, a proxy for composition, after magma ocean crystallization is shown by different colored layers. As overturn occurs, the dense upper layers sink and the less dense lower layers rise. The final product is a complicated mixture of materials.

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PSRD:A Primordial and Complicated Ocean of Magma on Mars

+ View Movie 2 In this movie we can track the initial location of a layer and see where it ends up in the complicated Martian mantle. As shown in Video 1 of initial density layering, the final product again is complicated. Another result of the calculations of crystallization and overturn is graphs of the concentrations of the main oxides that make up Martian crustal rocks. These show significant variations in oxide concentrations with depth. After overturn, these concentrations are completely different. In the middle mantle they oscillate back and forth, indicating significant scrambling of the mantle, as shown in the movies above. There are two distinct regions with different Al2O3 concentrations that could subsequently melt to produce most of the Martian meteorites (low in Al2O3) and crustal rocks at the Pathfinder, Spirit, and Opportunity landing sites.

Compositional stratification resulting from crystallization of a Martian magma ocean (without formation of a pure garnet layer). Original concentration profiles are drastically altered and there is a mixed-up region in the middle. Partial melting of different regions of the mantle will produce several types of magma, giving rise to a variety of igneous rocks.

Elkins-Tanton and her coworkers and Lars Borg and Dave Draper also investigated the distribution of trace elements in different regions of the mantle, and mathematically melted the mantle to produce magmas. The results are consistent with the picture being painted by Martian meteorites and surface rocks. These sketches http://www.psrd.hawaii.edu/Mar06/mars_magmaOcean.html

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PSRD:A Primordial and Complicated Ocean of Magma on Mars show that there is a wide variety of mantle compositional regions and that they formed early in the history of Mars. Elkins-Tanton suggests that the Shergottite group of Martian meteorites could have formed from a shallow region of the mantle that is depleted in Al2O3, whereas igneous rocks analyzed at the Pathfinder and Mars Exploration Rover sites formed from deeper regions of the mantle richer in Al2O3. It makes a nice, consistent story.

Spin-off: Driving the Early Martian Magnetic Field

Models of mantle overturn explain how we can get the diversity of igneous rocks observed on Mars. But wait, there's more. The models may also explain what created an early magnetic field on Mars. Magnetic measurements by a magnetometer onboard the Mars Global Surveyor spacecraft showed that there are regions of the ancient, heavily cratered highlands of Mars that have significant magnetic anomalies--areas with stronger than average magnetic fields. The fields are recorded by magnetic minerals in rocks.

Map of the magnetic field strength on Mars. Red colors indicate strong fields, other colors weaker fields. The strong fields occur in the Martian highlands--the oldest exposed parts of the crust.

Most geophysicists think that planetary magnetic fields are generated by convective motions inside metallic iron cores. The motions are driven by cooling of the core, which requires flow of heat into the mantle immediately above the core. Some of the mantle rock arrived as relatively cool cumulates from the upper parts of the magma ocean. This allows for significant flow of heat across the core-mantle boundary, driving convection in the core. However, the heat from the core raises the temperature of the mantle rock, as does heat released by decay of radioactive elements trapped in the cumulates. This slows down the heat flow, inhibiting and then stopping the core dynamo. Elkins-Tanton and her colleagues estimate that this takes between 15 and 150 million years. After that the magnetic field is much weaker. The areas with the magnetic anomalies, therefore, are the oldest on Mars. Other areas in the highlands either formed later or were modified significantly by intrusion of younger magma.

What Next?

The complex models developed by these two teams of cosmochemists and geophysicists make predictions about the types of magma produced on Mars. Continued searching for Martian meteorites in hot and cold deserts on Earth and continued analysis of igneous rocks on the Martian surface by orbital and landed spacecraft will allow tests of these predictions. If unpredicted rocks are found, the models can be modified and we will have learned more about magma ocean crystallization and overturn, and about subsequent partial melting and magma formation.

http://www.psrd.hawaii.edu/Mar06/mars_magmaOcean.html

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The Mars Science Laboratory rover will be much larger than the Mars Exploration Rovers Spirit and Opportunity. It will be equipped with instruments that will make analyses of surface rocks and soils, including igneous rocks that can be used to test predictions made by geochemical-geophysical models of formation of the Martian mantle.

2005-2006 ANSMET team members, Marie Keiding (Univeristy of Iceland), Gordon (Oz) Osinski (Canadian Space Agency), and Shaun Norman (ANSMET Mountaineer and Field Safety Leader) are shown collecting a meteorite from the Miller Range icefields. The continued search for meteorites in the cold desert of Antarctica and in hot deserts of northern Africa and elsewhere will continue to yield new Martian meteorites. Cosmochemists hope that new types are found, expanding our knowledge of the range of igneous rocks on Mars.

LINKS OPEN IN A NEW WINDOW.

z

z

z

Borg, L. E. and D. S. Draper (2003) A petrogenetic model for the origin and compositional variation of the martian basaltic meteorites. Meteoritics and Planetary Science, v. 38, p. 1713-1731. Elkins-Tanton, L. T., E. M. Parmentier, and P. C. Hess (2003) Magma ocean fractional crystallization and cumulate overturn in terrestrial planets: Implications for Mars. Meteoritics and Planetary Science, v. 38, p. 1753-1771. Elkins-Tanton, L. T., E. M. Parmentier, and P. C. Hess (2005) Possible formation of ancient crust on

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z

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Mars through magma ocean processes. Journal of Geophysical Research, v. 110, E12S01, doi:10.1029/2005JE002480. Elkins-Tanton, L. T., S. E. Zaranek, E. M. Parmentier, and P. C. Hess (2005) Early magnetic field and magmatic activity on Mars from magma ocean cumulate overturn. Earth and Planetary Science Letters, v. 236, p. 1-12.

Taylor, G. J. (1997) Moonbeams and Elements. Planetary Science Research Discoveries.

http://www.psrd.hawaii.edu/Oct97/MoonFeO.html Taylor, G. J. (2004) The Multifarious Martian Mantle. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/June04/martianMantle.html Taylor, G. J. (2005) Magma and Water on Mars. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/Dec05/Magma-WaterOnMars.html

Excellent recent reviews of Martian geological, geophysical, and geochemical evolution: z

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Nimmo, F. and K. Tanaka (2005) Early crustal evolution of Mars. Annual Reviews of Earth and Planetary Science, v. 33, p. 133-161. Solomon, S. and many others (2005) New perspectives on ancient Mars. Science, v. 307, p. 1214-1220.

[ About PSRD | Archive | Search | Subscribe ] [ Glossary | General Resources | Comments | Top of page ] 2006 [email protected] main URL is http://www.psrd.hawaii.edu/

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PSRD: Martian Meteorites Record Surface Temperatures on Mars

http://www.psrd.hawaii.edu/July05/Mars_paleotemp.html

posted July 29, 2005

Martian Meteorites Record Surface Temperatures on Mars --- Gases trapped in Martian meteorites indicate that Mars has been a cold desert for a long, long time.

Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology

Using published data for argon (Ar) released when Martian meteorites are heated, David Shuster (California Institute of Technology, now at Berkeley Geochronology Center, Berkeley, CA) and Benjamin Weiss (Massachusetts Institute of Technology) show that the nakhlite group of Martian meteorites and unique Martian meteorite ALH 84001 were probably not heated above about 0 oC for most of their histories. This indicates that the surface of Mars has been cold for almost four billion years. If a warm, wet environment existed on Mars (inferred from previous studies of surface features and geochemical parameters), it occurred before four billion years ago. Reference: Shuster, David L. and Benjamin P. Weiss (2005) Martian surface paleotemperatures from thermochronology of Meteorites. Science, vol. 309, p. 594-597.

Soaking Wet, Bone Dry Mars

Climate change on Earth is often in the news. Climate specialists worry about swings in global temperatures of several degrees Celsius. This does not sound like much, but it is enough to cause ice ages sometimes and widespread shallow seas at other times. But those changes are nothing compared to what the planet Mars seems to have experienced. Mars is decorated with huge channels eroded when vast quantities of water flowed through them. Oceans may have existed in the northern plains. Valley networks decorate the planet's surface. Yet now it is a dry, cold place. The daily average temperature at the equator is an ultra-nippy 60 oC below zero. Its monotonous dry climate has been enlivened occasionally by water seeping from the sides of impact craters, and changes in the planet's tilt may have moved glaciers from the current poles to more equatorial regions, but basically it has been colder and drier than anyplace on Earth. Yet at some time in the past, probably billions of years ago, Mars was a much warmer and wetter place.

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http://www.psrd.hawaii.edu/July05/Mars_paleotemp.html

Mars flaunts strong evidence for vigorous water activity in the past (see images below from left to right), such as immense, water-carve outflow channels, valley networks, possibly an extensive northern ocean, and presence of layered deposits whose origin involved evaporation of salty water.

On the other hand (see images below), it appears today to be extremely dry, a vast desert shaped mostly by wind, except in a limited number of locales where water has recently formed gullies on the walls of impact craters.

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http://www.psrd.hawaii.edu/July05/Mars_paleotemp.html

David Shuster and Benjamin Weiss wanted to determine past temperatures during this impressively long Martian cold, dry spell. Experts in determining the ages of rocks using potassium-argon dating and its advanced cousin, 40Ar/39Ar dating, they reckoned that Martian meteorites contained a record of surface temperatures. This is possible because Ar leaks out of rocks unless they are kept cool enough. They chose to study the nakhlite group of Martian meteorites because they do not have the same level of shock damage by meteoroid impact as do other types of Martian meteorites, thereby minimizing one form of heating besides surface temperature. They also studied data from Allan Hills (ALH) 84001 because it is by far the oldest in our collection of Martian meteorites. (For evidence that Martian meteorites actually do come from Mars, go to the curatorial office at the Johnson Space Center.) Nakhlites have already proven to be useful in assessing the timing of relatively recent aqueous events on Mars (see PSRD article: Liquid Water on Mars: The Story from Meteorites). The nakhlites contain mineral grains formed by the reaction of water with original minerals and deposition of others as the solutions dried up (see images below). Tim Swindle and his colleagues at the University of Arizona determined from potassium-argon dating that these water-based alteration events were of short duration and took place intermittently during the past 600 million years or so. Shuster and Weiss hoped to look at a broader time scale and to set limits on the temperature during the past 4 billion years.

The nakhlite group of Martian meteorites show that small amounts of water have flowed on Mars since the nakhlites formed in lava flows 1.3 billion years ago. On the left is a transmitted light photograph of red staining in an olivine crystal in the MIL 03346 nakhlite. The staining is composed of a complex mixture of weathering products. On the right is another transmitted light photograph of the same meteorite showing sulfate crystals deposited from evaporating salty water.

Using Ages to Deduce Temperatures

The startling thing about nakhlites is that all age dating techniques give the same age for their origin as igneous rocks, 1.3 billion years. You'd think that the concordance of ages by potassium-argon, rubidium-strontium, uranium-lead, and samarium-neodymium would not be surprising. If these techniques really work to date a rock and the rock formed at a given time, shouldn't they all yield the same result? They would, if nothing happened to a rock after it was formed, but it can be heated, metamorphosed, shocked (by impact), and altered by water. Because each age-dating technique is affected by these events differently, they tend to yield different apparent ages. The fact that the nakhlites give the same age for all systems indicates that they have had a relatively simple history. Even uranium-thorium-helium dating, which is easy to alter because helium leaks out of minerals like sand through a sieve (even at low temperatures), gives about the same age (0.8 to 1.2 billion years). 3 of 6

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40Ar/39Ar

age dating involves irradiating a sample with fast neutrons in a nuclear reactor and then measuring the ratio as the sample is heated progressively from low temperature (about 250 oC) to higher temperature (up to 1200 oC). Because 39Ar is produced from nuclear reactions with 39K, the experiment measures both potassium and 40Ar at the same time. This progressive heating causes release of the gases from different minerals sequentially, providing information about the temperature-time history of the rock. This can be quantified by knowing the rates at which argon diffuses out of mineral grains and the sizes of the mineral grains. The result is that the nakhlites appear to have lost only 1% of the 40Ar produced by the decay of 40K since they formed 1.3 billion years ago. The problem is that Shuster and Weiss did not know how hot the nakhlites got or for how long. Nevertheless, they could test different intensities of heating events to produce a set of solutions that result in loss of only 1% of the 40Ar from the nakhlites. The calculations are shown in the diagram below. If the nakhlites were never heated after they formed, they would preserve their age if held at a temperature of about minus 15 oC. If heated for a period of 10 million years, they could have reached as high as about 90 oC if the heating happened soon after they formed, but much less if more recent. A 10-million-year heating event that occurred during the past one billion years would not have heated the nakhlites to more than about 20 oC. Longer duration heating events must have been much cooler than 20 oC, and most likely not much higher than zero oC (see graph below).

Maximum temperatures reached in long-duration heating events of the nakhlite lava flows on Mars. The curves are calculated from gas-release data from the Nakhla meteorite. They show the maximum temperature reached for a temperature increase lasting for 10, 100, 200, and 500 million years, and beginning at any point along the curve. The case for no heating (isothermal) is also shown. A sustained period of warm temperature is possible, but it is more likely that the nakhlites were not heated to more than about zero o C during the past billion years or so.

A similar analysis can be done for the ancient ALH 84001 meteorite. The calculations in that case indicate that in all likelihood that meteorite, which was shock-heated 3.9 billion years ago, was not since heated above 0 oC for longer than a million years. In fact, it is likely that it was never hotter than about 7 oC for more than a million years during the past 3.9 billion years (see graph below). This suggests that Mars has been mostly a dry desert for all that time.

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The ancient meteorite ALH 84001 formed about 4.5 billion years ago, but was shocked (probably when excavated from great depth) 3.9 billion years ago. Since then, it could not o have been heated to much more than zero C for more than a million years.

One worry is that the nakhlites and ALH 84001 could have been heated when a big impact launched them from Mars. Shuster and Weiss address this problem. Using the same techniques to calculate gas loss, and knowing from the cosmic ray exposure ages of nakhlites that they were launched 11 million years ago, Shuster and Weiss calculate that the meteorites could not have been heated to more than 350 oC for more than a few hours.

The Long Drought

If Shuster and Weiss' analysis is correct, the areas on Mars that were home to the nakhlites and ALH 84001 got neither warm nor wet for very long. Short increases in temperature and brief wet spells are certainly allowable, and even required by the presence of weathering products in the nakhlites. Mars appears to have been a desert for billions of years. This implies that if life arose on the Red Planet, it is likely to be hidden underground. Places with groundwater beneath a permanently frozen underground cryosphere may be teeming with life. Or not. We can search for this life by drilling deep into the crust, or by choosing the right spots to sample, such as the terminations of young gullies and other apparently youthful features shaped by flowing water. Shuster and Weiss also point out that the lack of heating of the nakhlites and ALH 84001 when they were blasted off Mars indicates how easily undamaged materials can be lifted off Mars and sent to Earth. The inner planets might not be biologically isolated from each other. Life on Mars (if there is life on Mars) might be related to life on Earth. We may all be one big, solar system family.

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LINKS OPEN IN A NEW WINDOW.

Mars meteorites comprehensive page from Ron Baalke, Jet Propulsion Lab. Martian meteorites from the American Museum of Natural History. MIL 03346: New Martian meteorite found in Antarctica, News release from Case Western Reserve University (2004). Shuster, David L. and Benjamin P. Weiss (2005) Martian surface paleotemperatures from thermochronology of Meteorites. Science, vol. 309, p. 594-597. Taylor, G. J. (2000) Liquid Water on Mars: The Story from Meteorites. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/May00/wetMars.html

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PSR Discoveries: Outflow Channels in Chryse Planitia, Mars

http://www.psrd.hawaii.edu/June01/MarsChryse.html

posted June 14, 2001

Outflow Channels May Make a Case for a Bygone Ocean on Mars Written by Linda M.V. Martel Hawai'i Institute of Geophysics and Planetology

High-resolution elevation data from the Mars Orbiter Laser Altimeter (MOLA) onboard the Mars Global Surveyor (MGS) spacecraft have been analyzed recently in Chryse Planitia to test the hypothesis that large outflow channels emptied into an ocean in this region of Mars. Researchers Mihail (Misha) Ivanov (Vernadsky Institute of Geochemistry and Analytical Chemistry) and James Head (Brown University) collected quantitative MOLA information on channel patterns, continuity, and elevations where those patterns change or disappear into the northern lowlands. Their recently published report describes how the channels end, or become more subtle, at elevations very close to a previously mapped geologic contact interpreted by some to represent a shoreline of an ancient ocean. Ivanov and Head hypothesize that the change in channel topography is consistent with flow of water from a river into a submarine environment with possible deposition of sediments by density currents deep into the North Polar basin. Reference: Ivanov, M. A. and J. W. Head, III (2001) Chryse Planitia, Mars: Topographic configuration, outflow channel continuity and sequence, and tests for hypothesized ancient bodies of water using Mars Orbiter Laser Altimeter (MOLA) data, Journal of Geophysical Research,vol. 106, p. 3275-3295.

Chryse Planitia

Chryse Planitia, chosen in the mid-1970s for the landing site of Viking 1, is a relatively flat, low, broad plain just north of the Martian equator. Because many of the largest Martian outflow channels converge here, Chryse Planitia is an ideal setting to study channel patterns and depositional environments. More importantly, researchers have noticed that the distinctive textures and teardrop-shaped islands inside the channels change and disappear near the margins of Chryse Planitia. These changes have led some people to hypothesize that debris-laden rivers may have emptied their loads into a lake at Chryse Planitia or into an ocean that occupied the northern lowlands. The actual timing of this bygone ocean is unknown, but may be Hesperian to Early Amazonian.

What the MOLA Data Reveal

Ivanov and Head cite five main lines of evidence from MOLA data that support the hypothesis that large outflow channels emptied into an existing standing body of water in the northern lowlands of Mars in Hesperian-Early Amazonian times. Chryse Planitia is not a locally closed basin but instead opens into the North Polar basin. MOLA data show that Chryse Planitia is not a closed basin, as thought previously, but rather a low area with a gentle regional slope to the north-northeast. [See areas colored blue in the map below.] If Chryse Planitia were a closed depression, then we'd expect water and sediments to be confined inside. The absence of a topographic low within Chryse Planitia, and lack of a barrier between the channel mouths and the northern lowlands, suggests that channels emptied and spread out beyond Chryse into the North Polar basin.

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http://www.psrd.hawaii.edu/June01/MarsChryse.html

The distinctive patterns of the six largest channels end at relatively the same elevation even though the channels are of different ages and are separated by hundreds of kilometers over a total lateral distance of more than 2500 kilometers. The six major outflow channels studied by Ivanov and Head were Kasei, Maja, Simud, Tiu, Ares, and Mawrth Valles. They determined the elevation of the most distal part of each channel at its contact with adjacent plains deposits and found the elevations to be quite close, all occurring within a range of 284 meters. (In contrast, the entire northern lowlands, from the edge of the cratered terrain to the bottom of the North Polar basin, has an elevation range of more than 3000 meters.) One explanation for this similarity in elevation where each channel ends (approximately -3742 meters, with a standard deviation of 153 meters) is that it represents a base level or shoreline where the channels joined a large body of water. It might be where rivers met the ocean. The Hesperian-aged channels end at elevations close to a previously mapped geologic contact interpreted as a shoreline of an ancient ocean. In 1989, Timothy Parker and colleagues at the Jet Propulsion Lab mapped boundary contacts between landforms in the northern lowland plains. They outlined two contacts that are generally parallel to the southern boundary of the northern lowlands, and interpreted them as two separate highstands of a now vanished ocean. The MOLA map shown below is centered on the north pole and shows topography of the northern hemisphere of Mars. Low elevations are shown in shades of blue while the black lines indicate positions of contact 1 and contact 2. Contact 2 has a range in elevation of 4700 meters with a mean elevation of -3760 meters. The dramatic variations in elevations could be due to geologic activity (for example, surface uplift or tilting, subsidence, erosion, sedimentation, lava flooding, etc.) that may have occurred after the shoreline formed. Contact 1 has a mean elevation of -1680 meters, and its elevation varies by almost 11 kilometers. Ivanov and Head report that the mean elevation of the termini of Chryse channels (-3742 meters) is within 18 meters of the mean elevation of Contact 2.

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PSR Discoveries: Outflow Channels in Chryse Planitia, Mars

http://www.psrd.hawaii.edu/June01/MarsChryse.html

Mars has no actual sea level. The elevation designated as zero, therefore, is defined by the mean Martian radius, 3,382.9 kilometers, and the average atmospheric pressure of 6.1 millibars (6.1 thousandths of the Earth's atmosphere). If you were standing on the martian surface and the center of the planet were 3,382.9 kilometers beneath your feet, then you would be standing at 0 kilometers elevation, shown on this MOLA map in yellow.

Some of the channels continue out of Chryse Planitia for hundreds of kilometers into the North Polar basin, but their patterns are subdued and very different once past their recognized termini. Subtle elongated depressions in the MOLA data continue along the trend of the mapped channels beyond the channel termini as shown in the MOLA map below. Black arrows indicate where the channels (Simud on the left and Tiu on the right) lose their distinct patterns. Ivanov and Head consider the loss of distinctive channel patterns as corresponding to a reduction in the erosive power of the channel flows.

The distinctive change in channel pattern is consistent with rapid loss of energy as when a river discharges into a shallow submarine environment. Ivanov and Head report that the lateral extent of the subdued channels compares favorably with environments on Earth where rivers enter marine environments, moving from subaerial (more erosional) to submarine (more depositional) settings and distributing sediment widely downslope as density currents.

Relationship of the Chryse Channels with the Hypothesized Ocean

Try to picture the floods. The crashing, tumultuous torrents were large and dirty, carrying rocks, ice, and sediments through the 3 of 5

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channels. Debris-laden floodwaters scoured the landscape, cutting through the underlying rock, and when they spread out, where did the floodwaters go? Did the water merely sink underground? Did it fill the North Polar basin to make an ocean? Did it enter an existing ocean? We can only guess what happened. We need more information on the actual volumes of water and sediments involved in the floods and the timing of the floods. Ivanov and Head refer to different estimates of channel volumes (from Michael Carr, Victor Baker and others) to try to give us a better picture. Carr's estimate of the water volume of a single large flood in a Martian outflow channel is 300,000 km3. This is enough water to flood the entire North American continent by 30 meters! It would take at least 46 of these floods to fill the Martian northern lowlands to the level of Contact 2. Other estimates suggest that each channel may have filled the North Polar basin in separate events, requiring significant water loss between floods and refilling to essentially the same level. Ivanov and Head favor the hypothesis that the channels flowed into a preexisting standing body of water whose margins were already near the level of Contact 2 and cite the striking similarity of elevations where the channels end and the proximity of these elevations to the mapped Contact 2.

Mars Ocean Still Being Debated

Whether or not Mars had large bodies of standing water remains an unanswered question and not all investigators support the notion of a vast northern ocean. Photogeologic mapping of the proposed shorelines by Michael Malin and Kenneth Edgett in 1999 using high resolution images of Mars taken with the Mars Orbiter Camera (MOC) showed no features they would attribute to the action of water in a coastal environment. Other researchers contend that ridge networks in the northern lowlands are indicative of tectonic processes related to the Tharsis volcanoes. Tectonic features in this area of Mars, however, are not inconsistent with the possible presence of an ocean. Earth's ocean basins are prime examples of tectonic features. Confirming the presence of large bodies of standing water in Martian history will require a multifaceted approach. We'll need laser altimiters (MOLA) for topographic data, cameras for photogeologic mapping, infrared spectrometers (such as onboard 2001 Mars Odyssey) for mapping the distribution of minerals on the surface, gamma ray spectrometers (also on Odyssey) for mapping the surface distribution and abundance of chemical elements, as well as mineral and chemical studies of meteorites and rocks returned from Mars (to check for the presence of salts, for example.) If there were standing bodies of water on Mars billions of years ago they may have influenced the planet's atmosphere and climate, geology, environmental chemistry, and ultimately its capacity to support the emergence of life. These are the reasons why scientists, including Ivanov and Head, seek evidence of ancient oceans on Mars. The search may quench our collective thirst for knowledge about the Red Planet.

Baker, V. R., R. G. Strom, V. C. Gulick, J. S. Kargel, G. Komatsu, and V. S. Kale (1991) Ancient Oceans, Ice Sheets and the Hydrological Cycle on Mars, Nature, vol. 352, p. 589-594. Carr, M. H. (1996) Water on Mars, Oxford University Press, New York. Head, J. W., III, H. Hiesinger, M. A. Ivanov, M. A. Kreslavsky, S. Pratt, and B. J. Thomson (1999) Possible Ancient Oceans on Mars: Evidence from Mars Orbiter Laser Altimeter Data, Science, vol. 286, p. 2134-2137. Head, J. W., III, M. Kreslavsky, J. Hiesinger, M. Ivanov, S. Pratt, and M. Seibert (1998) Oceans in the Past History of Mars: Tests for Their Presence Using Mars Orbiter Laser Altimeter (MOLA) Data, Geophysical Research Letters, vol. 25, p. 4401-4404. Ivanov, M. A. and J. W. Head, III (2001) Chryse Planitia, Mars: Topographic configuration, outflow channel continuity and sequence, and tests for hypothesized ancient bodies of water using Mars Orbiter Laser Altimeter (MOLA) data, Journal of Geophysical Research,vol. 106, p. 3275-3295. Malin, M. C. and K. S. Edgett (1999) Oceans or Seas in the Martian Northern Lowlands: High Resolution Imaging Tests of Proposed Coastlines, Geophysical Research Letters, vol.26, p. 3049-3052. Mars Exobiology Strategy Mars Global Surveyor 4 of 5

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2001 Mars Odyssey MOLA Science Investigation. Parker, T. J., R. S. Saunders, and D. M. Schneeberger (1989) Transitional Morphology in West Deuteronilus Mensae, Mars: Implications for Modification of the Lowland/Upland Boundary, Icarus, vol. 82, p. 111-145. Smith, D. E., and others (1999) The Global Topography of Mars and Implications for Surface Evolution, Science, vol. 284, p. 1495-1503. Withers, P. and G. A. Neumann (2001) Enigmatic northern plains of Mars, Nature, vol. 410, p. 651. (University of Arizona press release.)

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PSRD: Dirty Ice on Mars

http://www.psrd.hawaii.edu/June02/MarsGRSice.html

posted June 5, 2002

Dirty Ice on Mars --- Instruments on the Odyssey spacecraft show that a lot of dirty ice sits within a meter of the surface in the south polar latitudes of Mars.

Written by G. Jeffrey Taylor Hawai'i Institute of Geophysics and Planetology

The gamma-ray and neutron spectrometers onboard the orbiting 2001 Mars Odyssey spacecraft have detected strong signals from hydrogen quite close to the Martian surface. The concentration of hydrogen is so large that it must be in the form of ice. The amount of ice in the upper meter or so begins to rise at about -60o latitude and continues to increase toward the South Pole. Detailed analysis of the data indicates that the ice-rich layer resides beneath a hydrogen-poor upper layer. The thickness of the upper layer decreases from about 75 cm at -42o to about 20 cm at -77o. The amount of ice in the lower layer is between 20 and 50 wt% (weight percent), with a best estimate of 35 wt%. Because ice is much less dense than mineral grains, this translates to more ice than rock by volume. It's dirty ice. The results were reported in papers by William Boynton (University of Arizona) and the gamma-ray team, by William Feldman (Los Alamos National Laboratory) and others, and Igor Mitrofanov (Russian Space Research Institute) and others. References: Boynton, W. V. and others (2002) Distribution of hydrogen in the near-surface of Mars: Evidence for subsurface ice deposits. Published online May 30 2002; 10.1126/science.1073722 (Science Express Reports.) Author list Feldman, W. C. and others (2002) Global distribution of neutrons from Mars: Results from Mars Odyssey. Published online May 30 2002; 10.1126/science.1073541 (Science Express Reports.) Author list Mitrofanov, I. and others (2002) Maps of subsurface hydrogen from the high-energy neutron detector, Mars Odyssey. Published online May 30 2002; 10.1126/science.1073616 (Science Express Reports.) Author list

Where Has All the Water Gone?

The surface of Mars has been sculpted by water. There are vast networks of valleys carved by flowing water (see example at left in S. Cerberus, 206o W, 8o N; click image for higher resolution options.) Immense channels were scoured by water gushing at hundreds of millions to a billion cubic meters per second. Some channels flow into craters and other depressions, forming smooth, flat deposits interpreted as former lakes. [See PSRD article: For a Cup of Water on Mars]. Some lava flows show strong evidence of interaction with subsurface ice [See PSRD article: If Lava Mingled with Ground Ice on Mars]. The striking fluidized ejecta surrounding many Martian impact craters may indicate that the target contained ice or water that volatilized during the impact, creating a runny debris flow that surged outwards from the growing crater. Some gullies apparently formed by wet debris flowing down steep slopes. Because the debris covers wind-blown deposits, they might have formed quite recently. An ocean might even have covered the northern plains of Mars. [See PSRD article: Outflow Channels May Make a Case for a Bygone Ocean on Mars].

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PSRD: Dirty Ice on Mars

http://www.psrd.hawaii.edu/June02/MarsGRSice.html

All this evidence points to a leading role for water in carving the Martian surface. But there are no rivers, lakes, oceans, or even puddles now. Where did all that water go? Most studies suggest much of it is deep underground. Some is trapped in mineral grains. For example, the Viking landers indicated that the soil contains about 1% chemically bound water. Estimates of how much water is underground vary, ranging from an amount that could cover all of Mars to a depth of 10 meters to as much as 1.5 kilometers. Most expert guesses come in around a few hundred meters. We need some direct measurements on how much water there is underground, but there have not been any direct measurements of the amount of water in the subsurface--until now. The Odyssey gamma-ray spectrometer suite of instruments analyzed the amount of water in the upper meter or so.

Gamma-Ray and Neutron Eyes

Mars is continuously bombarded with cosmic rays, which are mostly high-energy protons. The protons interact with the surface, causing assorted nuclear reactions. The reactions produce neutrons, which collide with surrounding nuclei. The nuclei become excited, and emit gamma rays as they return to their original, humdrum state. Gamma rays are a form of electromagnetic radiation; they have the shortest wavelength and highest energy. The gamma rays are characteristic of specific nuclear interactions in the surface, so measuring their intensity and wavelength allow a measurement of the abundance of several elements. One of these is hydrogen, which has a prominent gamma ray emission at 2.223 million electron volts (a measure of the energy of the gamma ray). This can be measured from orbit with the Gamma-Ray Spectrometer (GRS for short). The neutrons start out with high energies, so they are called fast neutrons. As they interact with the nuclei of atoms in the surface the neutrons begin to slow down, reaching an intermediate range called epithermal neutrons. The slowing-down process is not too efficient because the neutrons bounce off many nuclei without losing much energy (hence speed). However, when neutrons interact with hydrogen nuclei, which are about the same mass as neutrons, they lose considerable energy, becoming thermal, or slow, neutrons. (The thermal neutrons can be captured by other atomic nuclei, which then can emit additional gamma rays.) The more hydrogen there is in the surface, the more thermal neutrons relative to epithermal neutrons. Many neutrons escape from the surface, flying up into space where they can be detected by the neutron detector on Mars Odyssey. The same technique was used to identify hydrogen enrichments, interpreted as water ice, in the polar regions of the Moon.

A Huge Amount of Hydrogen

The Odyssey GRS did not detect much hydrogen in equatorial regions during the initial two months of measurements. It's winter in the northern hemisphere of Mars, so the ground in high northern latitudes is still blanketed by carbon dioxide frost, which obscures the signal from hydrogen. The southern hemisphere tells a different story. As we move from about -45o the hydrogen signal from the GRS increases steadily toward the South Pole. This is accompanied by a steady decrease in epithermal neutrons, which is also consistent with the presence of hydrogen. The GRS team had to convert the count rates of gamma rays and neutrons to the abundance of hydrogen. This is not a simple task because the relationship between concentration and gamma-ray signal is complex. It depends on whether hydrogen is distributed uniformly in the upper meter or so. To figure this out the U.S. members of the team used a computer program developed at Los Alamos National Laboratory to calculate the expected neutron and gamma-ray signals from surfaces with one or two layers with a variety of water (hence hydrogen) contents. This required some assumptions. A central one is that the bulk chemical composition of the soil is the same as the dirt measured at the Mars Pathfinder landing site. It also requires calibrating the measurements. That is not completely done as yet, so the GRS team normalized the data to 1 wt% H2O at the Viking landing sites, which are in equatorial regions. The calculations show that if the Martian surface layer, called the regolith, has H2O uniformly distributed, the flux of thermal neutrons escaping from the surface increases as epithermal neutrons decrease. This inverse correlation holds until H2O reaches about 10 wt%, at which point the thermal neutron flux decreases as the epithermal neutron flux decreases. The switch in behavior happens because hydrogen is such a strong moderator of neutron energy that when there's too much of it lots of the thermal neutrons are captured. If the regolith is layered, with a hydrogen-poor layer overlying one rich in hydrogen, the effect on the neutron flux is quite 2 of 7

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different. For a given water content in the upper and lower layers, as the thickness of the upper layer decreases, both the epithermal and thermal neutron fluxes decrease, until the upper layer becomes quite thin, less than about 20 cm. At that point, the thermal neutron flux increases rapidly as the epithermal flux slowly decreases. The Odyssey data show a correlated decrease in both epithermal neutrons and thermal neutrons, suggesting that the regolith is layered.

Calculated flux of epithermal vs. thermal neutrons for two cases. In one (right hand curve) the regolith is homogeneous, with the amount of H2O indicated on the curve. In the other (left curve) case the regolith is layered, with an upper layer containing 1 wt% water and the lower layer containing 35 wt%. The numbers along the curve represent the thickness of the upper layer, expressed in grams per square centimeters, which is roughly the same as the depth in centimeters if the density of the regolith is 1 g/cm3.

Physicists like to express thickness in grams per square centimeter, g/cm2. That's because they study the flux coming out from each square centimeter of the surface. To convert to depth, you divide by the regolith density. For example, suppose the depth is 60 g/cm2. A reasonable density of the regolith, based on measurements by the Viking landers, is about 1 to 1.3 g/cm3 (grams per cubic centimeter), which is mineral and rock grains with a lot of void space. Dividing 60 g/cm2 by 1 g/cm3 gives a depth of 60 cm. So, for the Odyssey results, as a rough rule of thumb, the depth in g/cm2 is the same as the depth in centimeters. The GRS team calculated the variation of epithermal and thermal neutrons for several cases. They used water contents in the upper layer of 1 or 2 wt% and water contents in the lower layer of 10, 20, 35, and 50 wt%. They also made the calculations for the thickness of the upper layer of 10 to 200 g/cm2. They compared the curves produced to the data obtained by the GRS onboard the Odyssey spacecraft. The data (red squares in the diagrams below) vary in the way we expect them to for a layered surface. The data for high latitudes (-62o to -77o) match the calculations for a lower layer containing 35 wt% H2O. The match is not as good for higher latitudes for the case shown below (2 wt% H2O in the upper layer), but matches much better if the upper layer contains only 1 wt% H2O (not shown in the diagrams below). This suggests that the amount of H2O in the upper layer might increase with latitude.

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Move your cursor over the buttons below to view curves that represent calculated neutron fluxes with the same H2O content in the lower layer (10%, 20%, 35%, or 50%) of a two-layer model assuming 2% H 2 O in the upper layer.

The dashed lines in the diagram below connect points on the curves corresponding to the indicated thickness of the upper layer. The correspondence with the measured data points between -62o and -77o suggests that the upper layer decreases in thickness toward the South Pole. A similar analysis can be made using the thermal neutron flux and the intensity of the hydrogen gamma-ray line. This also suggests that the H2O content of the lower layer is about 35 wt%.

Calculated thermal and epithermal neutron fluxes compared to data measured by the Odyssey instruments (red squares). The best match at high latitudes (-62 o and higher) occurs at 35 wt% H 2 O. Considering all uncertainties, the Odyssey GRS team believes that H2O is in the range 20 to 50 wt %. 4 of 7

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The ice-rich regions are found only where it is cold. This is consistent with calculations that predicted where ice should be stable beneath the Martian surface. The map of epithermal neutrons (below) shows where the flux is low, hence H2O is high. Measurements by the High Energy Neutron Detector (HEND) are consistent with this inference. The HEND insturments measured the distribution of epithermal and fast neutrons as well. It detected a deficiency of high-energy neutrons in the southern region., again indicating the presence of ice. The HEND team in Russia independently studied the structure of the near-surface layers and also concluded that it probably consists of an upper layer of regolith with about 5 wt% H 2O that covers a lower layer with a much higher percentage of H2O. The icy areas correspond to places where Michael Mellon and Bruce Jakosky predict that ice occurs within 80 cm of the surface (white lines on the map below).

Map of the epithermal neutron flux measured by the Neutron Spectrometer on Mars Odyssey (click for higher resolution option.) Low levels indicate the presence of high hydrogen concentration, interpreted as H 2 O or (especially for areas with higher epithermal neutron flux) as H 2 O or OH - bound in minerals. Note the high concentration of hydrogen south of -60 o latitude. Theoretical calculations predict that these regions would have ice within 80 cm of the surface.

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Map of the epithermal neutron flux measured by the High Energy Neutron Detector on Mars Odyssey (click for higher resolution option.) Low levels indicate the presence of high hydrogen concentration, interpreted as H 2 O or (especially for areas with higher epithermal neutron flux) as H 2 O or OH - bound in minerals. Note the very good consistency of independent measurements made by the Neutron Spectrometer (map above) and HEND instruments. Theoretical calculations made by the HEND team predict that southern regions would have ice below 60 cm of the surface.

The correspondence between the GRS data and the calculations is good, though not perfect. This suggests that the H2O is not distributed in two uniform layers with a sharp boundary between them. Nevertheless, the conclusion that the regolith beneath 20 to 40 cm contains 35 wt% H2O is completely consistent with the data. This corresponds to a lot of ice. Because the density of ice (0.9 g/cm3) is much less than that of mineral grains (about 2.5 g/cm3), 35 wt% ice corresponds to 60% ice by volume. It is dirty ice, not soil with ice in it. Some of the hydrogen might be sequestered in hydrous minerals, but most minerals do not contain such a large abundance of either H2O or other hydrogen molecules (such as OH-).

This diagram summarizes the interpretation made by the Odyssey GRS team. An upper, ice-poor layer overlies one rich in ice. Thickness of the ice-poor layer decreases to the south, to as low as about 20 centimeters. Even at -40 o , the ice-rich layer might not be too much more than a meter below the surface.

Climate History, Surface Processes, and Martian Settlers

As noted above, estimates of the amount of water on Mars center on a quantity equivalent to a global layer a few hundred meters thick. The Odyssey results cannot address how much water could be present beneath the one-meter depth in which the gamma rays and neutrons are produced. The whopping amount of water is certainly consistent with the idea that the upper kilometer or so 6 of 7

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of the Martian regolith might be a substantial reservoir of Martian water. The astonishing new discovery lies in the amount of ice and in the inferred pore space in the regolith. Pore space is simply the amount of empty space between solid, rocky grains. The pore space might have been produced by condensation of carbon dioxide frost in the winter. The condensing CO2 would contain tiny grains of dust. When the CO2 evaporated in the spring, it would leave behind a fluffy deposit of dust. Over time this fluffy, highly-porous deposit might occupy the upper meter of the regolith, providing a medium in which atmospheric water could condense as ice. If the water has been deposited by vapor exchange with the atmosphere, it holds clues to Martian climate. It may allow us to understand the atmospheric water cycle on Mars and perhaps even the history of the climate. This science story is just beginning. Understanding the climate history of Mars and the details of the deposition of ice in the upper meter will require more observations from Odyssey and future missions, but we do know something very important right now: Ice is very close to the surface in some places on Mars. This has enormous implications for human exploration and settlement of the planet. Although water can be extracted from hydrated minerals anywhere on Mars, it is much easier and cheaper to dig up some dirty ice, melt it, filter it, and use it. Water is essential for human life, for use in agriculture, and for converting into hydrogen and oxygen for fuel. There is a vast supply of water sitting within a meter of the surface for use by Martian settlers. The surprising results reported by the GRS team show how our knowledge leaps when we make measurements in new ways. Orbital remote sensing measurements of Mars have been made in visible, near infrared, and thermal infrared wavelengths. Extending the measurements to gamma rays and neutrons has opened up new vistas for understanding the geological history of Mars.

2001 Mars Odyssey homepage. 2001 Mars Odyssey Gamma Ray Spectrometer homepage. Boynton, W. V. and others (2002) Distribution of hydrogen in the near-surface of Mars: Evidence for subsurface ice deposits. Published online May 30 2002; 10.1126/science.1073722 (Science Express Reports.) Author list Feldman, W. C. and others (2002) Global distribution of neutrons from Mars: Results from Mars Odyssey. Published online May 30 2002; 10.1126/science.1073541 (Science Express Reports.) Author list Mellon, M. T. and Jakosky, B. M. (1993) Geographic variations in the thermal and diffusive stability of ground ice on Mars. Jour. Geophys. Res., vol. 98, p. 3345-3364. Mitrofanov, I. and others (2002) Maps of subsurface hydrogen from the high-energy neutron detector, Mars Odyssey. Published online May 30 2002; 10.1126/science.1073616 (Science Express Reports.) Author list

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PSRD: Did Martian Meteorites Come From These Sources?

http://www.psrd.hawaii.edu/Jan07/MarsRayedCraters.html

posted January xx, 2007

Did Martian Meteorites Come From These Sources? --- Researchers find large rayed craters on Mars and consider the reasons why they may be launching sites of Martian meteorites.

Written by Linda M. V. Martel Hawai'i Institute of Geophysics and Planetology

Large rayed craters on Mars, not immediately obvious in visible light, have been identified in thermal infrared data obtained from the Thermal Emission Imaging System (THEMIS) onboard Mars Odyssey. Livio Tornabene (previously at the University of Tennessee, Knoxville and now at the University of Arizona, Tucson) and colleagues have mapped rayed craters primarily within young (Amazonian) volcanic plains in or near Elysium Planitia. They found that rays consist of numerous chains of secondary craters, their overlapping ejecta, and possibly primary ejecta from the source crater. Their work also suggests rayed craters may have formed preferentially in volatile-rich targets by oblique impacts. The physical details of the rayed craters and the target surfaces combined with current models of Martian meteorite delivery and cosmochemical analyses of Martian meteorites lead Tornabene and coauthors to conclude that these large rayed craters are plausible source regions for Martian meteorites. Reference: Tornabene, L. L., J. E. Moersch, H. Y. McSween Jr., A. S. McEwen, J. L. Piatek, K. A. Milam, and P. R. Christensen (2006) Identification of large (2-10 km) rayed craters on Mars in THEMIS thermal infrared images: Implications for possible Martian meteorite source regions. Journal of Geophysical Res., v. 111, doi: 10.1029/2005JE002600).

Finding What They're Looking For

There are currently 34 Martian meteorites identified out of the 24,000+ that have been cataloged. The numbers are growing as a result of ongoing searches primarily in the world's deserts (for example see PSRD article: Searching Antarctic Ice for Meteorites). Cosmochemists have determined that these rocks came from basaltic igneous sources with young (by planetary standards) crystallization ages no more than 1.3 billion years (with the one exception: ALH84001 with an age of 4.5 billion years) and were ejected from Mars by impact cratering events between 600,000 and 20 million years ago. While these rocks provide invaluable direct 'ground truth' that scientists are using to help piece together the chemical and geological history of Mars, the question remains where exactly did these rocks come from? Knowing their provenance will add significant details to our understanding of how the planet formed, differentiated, and evolved geologically. One approach to answering the question has been to search orbital multispectral datasets to find volcanic terrains on Mars that match the mineralogy and spectral properties of Martian meteorites. These locales must be sufficiently dust-free to allow spectral analysis of the surface compositions and must also have at least one impact crater of appropriate size and age that could have ejected rocks at greater than Mars' escape velocity of ~ 5 kilometers per second. Previous work by Vicky Hamilton (University of Hawaii) using data from the Mars Global Surveyor Thermal Emission Spectrometer (TES) pointed to Eos Chasma, a branch of the Valles Marineris canyon system, as a possible source for unique Martian meteorite ALH84001. Hamilton's work with Ralph Harvey (Case Western Reserve University) identified Syrtis Major as a possible source region of nakhlite/chassignite meteorites. This is exciting on-going work to find meteorite source regions.

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Alternatively, an answer to Martian meteorite sources may well come from studies of some uncommon Martian craters that, until just a few years ago, had gone unnoticed. In 2003, using new Mars Odyssey Thermal Emission Imaging System (THEMIS) thermal infrared data, Alfred McEwen (University of Arizona) and colleagues reported the first discovery of a rayed crater, Zunil, in the southern Elysium region of Mars. More recently, Livio Tornabene and colleagues have identified an additional four large rayed craters and three more they deem probable. Their detailed observations of the craters combined with the known geochemistry of the meteorites and models of how meteorites are ejected off the planet add up to a compelling story that these rayed craters could have supplied Martian meteorites.

Rayed Craters Defined and Located

Tornabene and coauthors define a crater ray as filamentous (thread-like) elements in radial to subradial lineaments that spread out from a source crater like spokes from the center of a wheel. A ray contrasts with the surrounding, underlying surface. We are used to seeing crater rays on the Moon in visible-light images where this contrast is recognized as albedo. Lunar rays are brighter than the underlying surface. On Mars crater rays are not distinctive in visible light but are apparent in the thermal infrared because of a thermophysical (temperature-related) contrast with the surroundings. Martian rays appear brighter or darker than the underlying surface depending on the relative thermal properties of the materials when (day or night) the images were taken (see images below). Moon

Mars

LEFT: Clementine images show albedo (reflectivity) variations on the nearside of Earth's Moon. Extensive bright ray systems surround craters Copernicus (upper left center) and Tycho (near bottom). Click the image for more information in a new window. Bright crater rays have also been observed on Mercury and the icy Galilean moons. RIGHT: Crater rays on Mars show up in the thermal infrared and can be dark or bright. This is a nighttime thermal infrared image of Gratteri crater and its dark rays. Rocky areas are brighter because they retained daytime warmth into the night. In contrast, the dark rays and patches show where finer-grain materials cooled after local sunset. Click the image for more information in a new window.

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Pairs of images illustrate how Martian crater rays appear in different wavelengths of light. The left hand images are THEMIS nighttime thermal infrared images. TOP LEFT (a): Gratteri crater rays are revealed as dark streaks (see figure caption above for more explanation). BOTTOM LEFT (c): Warmer (brighter) Gratteri ejecta around the crater and colder (darker) rays are revealed. The right hand images are contrast-stretched MOC visible images of the same area. RIGHT (b and d): Crater rays are not easily seen. There is very little albedo difference between rays and the surrounding plains.

In their global survey for rayed craters, Tornabene and colleagues studied both THEMIS nighttime and daytime thermal infrared (TIR) brightness temperature images derived from band 9 (wavelength of 12.57 micrometers). This wavelength is used because it has the highest signal to noise value and is relatively transparent to atmospheric dust. Image resolutions ranged from 32 to 256 pixels per degree. They found that daytime TIR images are not as useful as nighttime TIR images for identifying rays. The effects of albedo, surface slope, and time of day all affect daytime surface temperatures more than at night. Moreover, the survey focused on latitudes specifically between 45 oN and 45 oS (see map below) because surface radiance and diurnal (a single daily cycle) thermal contrast generally decrease poleward of the equator. They note that lower surface radiance translates as lower signal to noise detected at the spacecraft's instrument, which generates poor quality images (especially at night when temperatures are much colder) making it more difficult to recognize rays with certainty at the higher latitudes. Their survey resulted in an additional four, and another three probable, rayed craters.

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The five known rayed craters (black circles with white name lables) and three probable rayed craters (white circles) are located on this MOLA shaded relief map of Mars. The rayed craters cluster in two groups; two craters are located south of Tharsis (the largest volcanic region on the planet) and the other six, including the probables, are in or near Elysium Planitia (the second largest volcanic region).

Characteristics of Martian Rayed Craters

Tornabene and colleagues found that the eight identified rayed craters all lie within lava plains or adjacent to the two major volcanic regions (Tharsis and Elysium Planitia, see map above). The table below (and continuing on the next page) summarizes what is known about Martian rayed craters. Each is shown in THEMIS nighttime infrared images, with North to the top. Crater locations are listed in latitude, longitude. Crater diameters and the longest ray length measured for each crater are listed in kilometers.

Martian rayed crater characteristics Zunil

Tomini

Crater Location: 7.7 oN, 166 oE southeast of Cerberus Planum within Elysium Planitia in Amazonian-aged lava plains Crater diameter: 10.1 km

Crater Location: 16.27 oN, 125.9 oE southwest of Elysium Mons in Hesperian-aged ridged volcanic plains Crater diameter: 7.4 km

Ray Length: 927 km Rays are not symmetrically arranged around the crater. White arrows point to two distinct dark rays to the northwest of Zunil crater.

Ray Length: 668 km Curved dark rays are prominent to the west and southeast. A long, straight ray trends north-northeast opposite the forbidden zone (a wedge-shaped zone where crater ejecta was never deposited).

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Gratteri

Zumba

Crater Location 17.7 oS, 199.9 oE near Memnonia Fossae southwest of Tharsis in Noachian-aged volcanic plateau Crater diameter: 6.9 km

Crater Location 28.65 oS, 226.9 oE Daedalia Planum in Hesperian-aged lava flows Crater diameter: 3.3 km

Ray Length: 595 km More than 30 rays, more than double the number found with any other rayed crater. Longest rays occur to the northwest and southeast. Region north of the crater is very dusty, so if rays exist in this region they lack a theromphysical contrast to the dust.

Ray Length: 240 km One of its longest rays trends to the east while another of the longest rays trends west. Zumba rays are unique because they are also distinct in daytime TIR images.

Dilly

Tomini B

Crater Location 13.27 oN, 157.23 oE Cerberus Planum within Elysium Planitia in Amazonian-aged volcanic terrain Crater diameter: 2.0 km with an elliptical shape

Crater Location 14.9 oN, 123.25 oE near Tomini crater in Hesperian-aged volcanic plains Crater diameter: 4.2 km

Ray Length: 50 km Dark rays appear to the northwest and southeast in distinctive "butterfly wing" pattern.

Ray Length: 220 km Three discernable, but faint rays.

Crater A

Crater B

Crater Location 18.1 oN, 155.5 oE near Zunil and Dilly craters in Amazonian-aged volcanic terrain Crater diameter: 5.7 km

Crater Location 15.5 oN, 159.2 oE near Zunil and Dilly craters in Amazonian-aged volcanic terrain Crater diameter:1.5 km

Ray Length: 86 km; Faint rays.

Ray Length: 42 km; Faint rays.

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Table Notes: The four new rayed craters have International Astronomical Union approved names of Gratteri, Tomini, Zumba, and Dilly named after small towns or villages in Italy, Indonesia, Ecuador, and Mali, respectively. The three crater names in italics are probable rayed craters. Probable rays were identified by Tornabene and coauthors after they used digital image processing to stretch the contrast on the nighttime thermal infrared images. The technique of contrast stretching reveals the subtle tonal differences in the images that represent faint temperature signatures of the rays. The white boxes in some of the images refer to close-up images in Tornabene and coauthor's published research article. The researchers found that Martian rays, like their lunar counterparts, are comprised of densely concentrated and clustered chains of small secondary impact craters (tens to hundreds of meters in diameter), ejecta, and surge deposits from the primary and secondary craters. Their detailed work shows that the rays have two characteristic thermal signatures: colder areas and warmer small spots. The different signatures are attributed to different particle sizes and induration (or hardness) of the materials. The colder areas (darker streaks in the THEMIS nighttime thermal infrared images shown in the table) are interpreted to be fine-grained, loosely-packed debris. This material cools quickly after sunset. On the other hand, the warmer (brighter) spots are interpreted to be fresh secondary craters (formed from ejecta blocks originating from the source crater), which excavated courser, rocky materials upon re-impacting the surface. Rocks or indurated sediments cool more slowly and consequently appear warmer (brighter) than their surroundings at night. A variety of other Mars data sets (such as THEMIS visual, Mars Orbiter Camera (MOC) narrow-angle images, TES-dervied thermal inertia, albedo, and dust cover maps) were compared to the THEMIS thermal infrared data to corroborate the interpretations made by Tornabene and his colleagues. The research team identified rayed craters in volcanic plains that are specifically characterized by intermediate values of albedo, thermal inertia, and dust cover index (previously defined as "thermophysical Unit C" by Michael Mellon, University of Colorado, Boulder, and colleagues). Only about 20% of the Martian surface appears to have this optimal combination of thermophysical properties needed to recognize crater rays. This suggests that other rayed craters may be present on Mars, but cannot be readily detected by means of THEMIS thermal infrared images. For example, if ray debris lies on top of surfaces that have dark nighttime TIR signatures, such as dust-mantled surfaces, then the rays are not discernable because there is no thermal contrast. Hence, a surface covered with a thick dust mantle obstructs our view of crater rays. Conversely, regions with little to no dust cover would not be able to produce the cold ray material that we so readily observe in THEMIS nighttime thermal infrared images. As a consequence, other (more rigorous and difficult) means may be necessary to detect additional rayed crater systems on Mars such as linking far-field secondary crater populations to a single source crater. High-resolution visible images from cameras like HiRISE (High Resolution Imaging Science Experiment) or CTX (Context Camera) on NASA's Mars Reconnaissance Orbiter may be very useful in future ray surveys.

Ray Patterns and Oblique Impacts

Rays are evidence of high-velocity ejecta. The fact that Martian rayed craters are among the freshest craters of their size and are found in young volcanic plains makes some people wonder if some of the ejecta from these impacts could have escaped from Mars to become Martian meteorites on Earth. The Mars rock would have to reach the escape velocity of 5 km/sec. During an impact event the kinetic energy of the incoming projectile causes shock deformation, heating, melting, and vaporization, as well as excavation of the crater and ejecta material. However, Martian meteorites show low to only moderate degrees of shock. To address this question Tornabene and coauthors examined the ray patterns to better understand the formation process. Elliptical crater shapes, forbidden zones (wedge-shaped zones lacking crater ejecta), and "butterfly wing" ray patterns in four definitive rayed craters (the exception is Zunil) and all three probables are cited as evidence that the Martian rayed craters formed by moderately oblique impact events. During oblique impacts it is possible that some of the ejecta is released at high velocities but low shock pressures. This is a process known as spallation and it is likely responsible for creating some of the ray-forming secondaries. Spallation is also currently the favored mechanism for ejecting relatively low-shocked rocks off Mars. Models also show increases in spallation volumes when oblique impacts strike volatile-rich subsurfaces. This is significant because Tornabene and colleagues observed fluidized ejecta blankets around the primary craters or around nearby larger craters--commonly recognized as evidence for subsurface volatiles (water ice). They concluded the ray formation process is consistent with spallation models of Martian meteorite delivery.

Launch Sites for

Cosmochemical analyses show that Martian meteorites came from basaltic igneous sources (by crystallization from cooling magma) with young (<1.3 billion year) crystallization ages (with the one exception of 4.5 billion year old ALH84001) and were blasted off Mars by impact cratering events 600,000 to 20 million years ago. Taking into account the uncertainties in ages derived by crater counting for Martian terrains (e.g. see review by William Hartmann, Planetary Science Institute, Tuscon), Tornabene and coauthors have suggested matches between certain rayed craters (based on surface ages) and the crystallization 6 of 7

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age groupings of the Martian meteorites. Based on their studies, Tornabene and colleagues suggest rayed craters within Elysium are possible sources for the shergottites and the two rayed craters outside of Elysium (Zumba and Gratteri) are possible sources for nakhlites, chassignite, and ALH 84001. Specifically, Zumba is in late Hesperian-age volcanic terrain and could be a source crater for the nakhlites and chassignites, which are about 1.3 billion years old. Gratteri, which is in older volcanic terrain (Noachian-age), is suggested as a possible source for the oldest Martian meteorite known, ALH 84001. As the research by Tornabene and coauthors shows, observations of the rayed craters and target surfaces combined with current models of Martian meteorite delivery and cosmochemical analyses of Martian meteorites suggest these large rayed craters are plausible source regions for Martian meteorites. Finding meteorite source regions will continue to pique our interest as researchers look further into the spectral signatures of the surfaces where rayed craters have been identified to help define and compare the compositions to the only field samples we have.

LINKS OPEN IN A NEW WINDOW.

Christensen, P. R., N. S. Gorelick, G. L. Mehall, and K. C. Murray, THEMIS Public Data Releases, Planetary Data System node, Arizona State University. http://themis-data.asu.edu. Hamilton, V. E. (2004) Detailed investigation of a globally unique, orthopyroxene-rich deposit in Eos Chasma, Mars, Eos Trans. AGU, 85, Fall Meet. Suppl., Abstract P11A-0959. Hartmann, W. K. (2005) Martian cratering 8: Isochron refinement and the chronology of Mars, Icarus, v. 174, p. 294-320. Harvey, R. P. and V. E. Hamilton (2005) Syrtis Major as the source region of the Nakhlite/Chassigny group of Martian meteorites: Implications for the geological history of Mars, Lunar Planet. Sci., XXXVI, Abstract #1019, Houston (CD-ROM). Mars Meteorites, comprehensive site from Ron Baalke, Jet Propulsion Lab.http://www2.jpl.nasa.gov/snc/. Martel, L. M. V. (2002) Searching Antarctic Ice for Meteorites. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/Feb02/meteoriteSearch.html. McEwen, A. S., E. Turtle, D. Burr, M. Milazzo, P. Lanagan, P. Christensen, J. Boyce, and the THEMIS Science Team (2003), Discovery of a large rayed crater on Mars: Implications for recent volcanic and fluvial activity and the origin of Martian Meteorites, abstract, 34th Lunar and Planetary Science Conference, Lunar Planet. Inst., Houston, TX. McEwen, A. S., B. S. Preblich, E. P. Turtle, N. A. Artemieva, M. P. Golombek, M. Hurst, R. L. Kirk, D. M. Burr, and P. R. Christensen (2005), The Rayed Crater Zunil and Interpretations of Small Impact Craters on Mars, Icarus, v. 176, p. 351? 381. McSween Jr., H. Y. (1999) Meteorites and Their Parent Planets. 2nd Edition. Cambridge University Press, 310 p. Mellon, M. T., B. M. Jakosky, H. H. Kieffer, P. R. Christensen (2000) High-reolution thermal inertia mapping from the Mars Global Surveyor Thermal Emission Spectrometer, Icarus, v. 148, p. 437-455. Taylor, G. J. (2005) Martian meteorites record surface temperatures on Mars. Planetary Science Research Discoveries. http://www.psrd.hawaii.edu/July05/Mars_paleotemp.html. THEMIS Feature Image: Gratteri Crater's Far-Flung Rays.

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2007 [email protected] main URL is http://www.psrd.hawaii.edu/

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PSRD: Recent Activity on Mars: Fire and Ice

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http://www.psrd.hawaii.edu/Jan05/MarsRecently.html

posted January 31, 2005

Recent Activity on Mars: Fire and Ice --- New images from Mars Express show evidence of recent volcanic and glacial activity on Mars, consistent with what we know from Martian meteorites and previous evaluations of the planet's internal heat production and climate.

Written by Linda M. V. Martel Hawai'i Institute of Geophysics and Planetology

Scientists combined the time-honored method of counting craters to estimate the age of planetary surfaces with brand new high-resolution, stereo images of Mars to reassess the planet's recent volcanic and glacial activity. Gerhard Neukum (Freie Universität, Berlin, Germany) and colleagues from Germany, United States, Russia, and the United Kingdom studied calderas on five major volcanoes and the shield of Olympus Mons with the High Resolution Stereo Camera (HRSC) on the European Space Agency's Mars Express Spacecraft to try to determine the duration of geologic activity more precisely than had ever been done before. Their work confirms that the Tharsis and Elysium regions were volcanically active over billions of years, that caldera eruptions were episodic but were especially numerous 100 to 200 million years ago, and that the most recent lava flows on Mars may be as young as two million years. Their findings are consistent with previous studies of Mars Global Surveyor data as well as Martian shergottite meteorites that suggest intermittent magmatism from 165 to about 500 million years ago. Neukum and coauthors also report the most recent phase of glacial activity on Olympus Mons was within the past four million years. So recent are these events in geologic time that the researchers speculate high-altitude, insulated ice deposits may be present on Olympus Mons even now and that volcanoes might still be active. Reference: Neukum, G., Jaumann, R., Hoffmann, H., Hauber, E., Head J. W., Basilevsky, A. T., Ivanov, B. A., Werner, S. C., van Gasselt, S., Murray, J. B., McCord, T., and the HRSC Co-investigator team (2004) Recent and episodic volcanic and glacial activity on Mars revealed by the High Resolution Stereo Camera. Nature, v. 432, p. 971-979.

New Ages

The High-Resolution Stereo Camera (HRSC) co-investigator team targeted the summit calderas of five major shield volcanoes and the flanks of Olympus Mons, known sites of relatively recent volcanic activity. They defined specific terrain areas for counting craters using the 10 meters/pixel resolution images and the Super Resolution Channel's 2.5 meters/pixel resolution data along with nested Mars Orbiter Camera (MOC) images. Previous age determinations using the crater counting technique have been limited by poorer resolution or by the small areas imaged. For a short explanation of the crater counting technique visit the Planetary Science

PSRD: Recent Activity on Mars: Fire and Ice

http://www.psrd.hawaii.edu/Jan05/MarsRecently.html

Institute web page.

Crater size-frequency distribution for the caldera on Arsia Mons. Frequencies of craters per unit area were counted on the caldera floor and plotted against crater diameters to derive the absolute age of the surface at 130 million years old. Similar log-log plots were constructed for each summit caldera.

The HRSC team used the unified cratering chronology model published in 2001 by Neukum and colleagues Boris Ivanov (Russian Academy of Sciences, Moscow) and William Hartmann (Planetary Science Institute, Tucson, Arizona) that concluded craters on the Moon and Mars were created by the same family of projectiles and that the lunar cratering chronology could be transferred to Mars. The ages derived from the Martian crater counts are limited in accuracy, however. The main uncertainty is the statistical error arising from the number of craters counted (the error increases as age, hence the number of craters, decreases). There is also uncertainty in the underlying impact flux model used for Mars relative to the lunar value. In the team's study, errors are approximately 20% to 30% for derived absolute ages younger than 3 billion years and 100 to 200 million years for ages older than 3 billion years. The images below show the five volcanic calderas examined in this study: Hecates Tholus, Albor Tholus, Arsia Mons, Ascraeus Mons, and Olympus Mons. In the left column are HRSC images of the five volcanic caldera complexes. They should appear as depressions (sun shining from the lower left corner). North is to the top of each image. In the right column are the crater-counting areas defined by Neukum and colleagues, which are labeled with their derived absolute ages.

Neukum and coauthors mapped five different caldera collapse events on Hectates Tholus with floors ranging in age from approximately 2 billion years old to about 100 milliion years old. [Additional images from ESA Mars Express.]

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PSRD: Recent Activity on Mars: Fire and Ice

http://www.psrd.hawaii.edu/Jan05/MarsRecently.html

The four caldera floors mapped on Albor Tholus range in age from approximately 2.2 billion years old to about 500 million years old.

The summit of Arsia Mons is dominated by a single huge caldera. The floor is dated at about 130 millions years old.

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PSRD: Recent Activity on Mars: Fire and Ice

http://www.psrd.hawaii.edu/Jan05/MarsRecently.html

A large central caldera floor, dated at approximately 100 million years old, cuts adjacent caldera floors of various ages.

Olympus Mons is unusual by comparison with the other caldera complexes because the five caldera segments have ages clustering around 100 to 200 million years old. In the simplest stratigraphic sense, the top-most features are the youngest. However, the absolute ages assigned to the floors of the five calderas on Olympus Mons seem to defy this and some would argue against the validity of the assigned ages. Nevertheless, the researchers explain that since the different ages are very similar within the error limits of ± 50 million years, we must keep in mind that the formation of all the calderas on Olympus Mons could have happened in a short time span around 150 million years ago. [Additional information and high resolution images from ESA Mars Express.]

The image below shows an area of the western scarp of Olympus Mons where sparsely cratered lava flows were reexamined with the HRSC and MOC images.

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PSRD: Recent Activity on Mars: Fire and Ice

http://www.psrd.hawaii.edu/Jan05/MarsRecently.html

Surfaces were dated using the crater size-frequency technique. The shaded areas show crater counting areas. Ages of the lava flows range from 115 million years to about two million years in the area where HRSC and MOC data were combined. [ Additional information and high resolution MOC image.]

Internal heat and magma supply

The calderas on five major volcanoes have undergone repeated activity as shown by the different ages of caldera floors created by different collapse events. Based on the crater size-frequency measurements by the HRSC team in these multiple calderas, magma reservoirs were forming, solidifying, and reforming on time scales of about 20 million years. The very long activity of Martian volcanoes implies correspondingly long lifetimes of hot spots in the planet's interior. These findings by Neukum and coauthors are in agreement with theoretical analyses and geological studies that suggest subsurface magma reservoirs must cool and solidify between caldera collapse events. Magma supply to the major shield volcanoes on Mars was episodic rather than continuous. What's more, the youngest volcanic surfaces in the study areas are so geologically young (about two million years) that volcanoes must have been active within the last 2% to 4% of Martian history. The HRSC team's conclusions for recent volcanic activity and an internally active Mars are completely consistent with what we know from other studies: Martian meteorites. Most of the Martian meteorites formed in lava flows or shallow magma bodies during Amazonian times. Shergottites are the most abundant type. All the basaltic shergottites have crystallization ages less than 500 million years, with eight of the rocks in the range of 165 to 180 million years. Although these individual rocks are far from a representative sampling of Martian lava flows, they do suggest intermittent magmatism during the past 500 million years and there is no reason to think there couldn't be even younger volcanic deposits.

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PSRD: Recent Activity on Mars: Fire and Ice

http://www.psrd.hawaii.edu/Jan05/MarsRecently.html

Previous interpretations of Mars Global Surveyor (MGS) camera, and Mars Orbiter Laser Altimeter (MOLA) data. Results from 2001 studies suggested sporadic volcanic activity lasted 100 million years or longer with estimated ages of 3 to 10 million years old for the youngest surfaces. Based on these young surface ages, studies of eruption rates, and the episodic eruption style, many researchers, including Susan Sakimoto (NASA Goddard Earth Science and Technology Center) concluded that the potential for volcanic eruptions occurring during the next several tens of millions of years on Mars was not out of the question. [See Leonard David's space.com article: Mars Volcanoes Still Alive After All These Years?.] Current interpretations of the Mars Odyssey Gamma Ray Spectrometer data and geophysical modeling. Using the crustal average values for potassium and thorium from Mars Odyssey GRS data, the GRS team has calculated that half of the planet's potassium and thorium is in the crust. That leaves half in the mantle. [Mars Odyssey GRS results are about to be submitted for publication, but see Taylor, G. J. et al. (2003) Igneous and aqueous processes on Mars: Evidence from measurements of K and Th by the Mars Odyssey Gamma Ray Spectrometer. Sixth International Conference on Mars, 3207 (pdf file).] Steve Hauck and Roger Phillips (Washington University, St. Louis) and Walter Kiefer (Lunar and Planetary Institute, Houston) calculated that if approximately half of the radioactive elements such as potassium, thorium, and uranium were left in the mantle, there would be enough heat source for a small amount of volcanic activity today on Mars.

The Time Scale is Geologic not Human

If the volcanoes on the Red Planet are potentially still active, then eruptions could occur. But when? Would any of the active ESA and NASA Mars missions record the event? The most reasonable forecast for any possible future volcanic activity is in another couple to tens of millions of years...well into the future by human standards. To put it into perspective, the lava flows on Mars are akin to those that erupt from volcanoes in Hawaii. But an eruption on Oahu (where the youngest volcanic rocks are only about 100,000 to 500,000 years old) is more likely than on Mars. Is there a chance of seeing a volcanic eruption on Oahu in our lifetime? A geologist will tell you it is possible, but the odds are actually slim. The same may be said for Mars.

Glaciers on the Shield of Olympus Mons

Glacial deposits at the base of the Olympus Mons scarp look like rock glaciers or debris-covered glaciers on Earth and are interpreted as evidence for repeated phases of activity. Rock glaciers are typically covered by rocks and boulders and often have ridges, furrows, and lobes on the surface. [See Milkovich and Head (2003) Olympus Mons Fan Shaped Deposit Morphology: Evidence for Debris Glaciers. Sixth International Conference on Mars, 3149 (pdf file) for a review of glacier types.] Neukum and coauthors found that the crater size-frequency distributions for these deposits ranged from 130 to 280 million years for the major lobes, 20 to 60 million years for some subareas, and four million years for the youngest surfaces. Snow/ice deposition on the Olympus Mons shield at elevations higher than 7000 meters may have led to episodes of glacial activity at this height. The data suggest that water ice protected by an insulating layer of dust may now exist at high altitudes at the edge of the Olympus Mons shield. Accumulations of water ice in non-polar regions are particularly hot topics of research because of their implications for hydrothermal activity and the strategy for searching for life on Mars.

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PSRD: Recent Activity on Mars: Fire and Ice

http://www.psrd.hawaii.edu/Jan05/MarsRecently.html

This perspective view of the western scarp of Olympus Mons shows steep gullied slopes, channels, and glacier-like flows.

More Mars Express Results

The first Mars Express Science Conference will be held February 21-25, 2005 in Noordwijk, The Netherlands. The scientific community involved in Mars Express will review the progress toward understanding Mars and put the results in the broader context of the latest scientific interpretations derived from current NASA missions: Mars Global Surveyor, Mars Odyssey, and Mars Exploration Rover.

LINKS OPEN IN A NEW WINDOW.

Borg, L. E., Nyquist, L. E., Weissman, H., Shih, C.-Y., and Reese, Y. (2003) The age of Dar al Gani 476 and the differentiation history of the Martian meteorites inferred from their radiogenic isotopic systematics. Geochim. Cosmochim. Acta, v. 67, p. 3519-3536. David, L. (2001) Mars Volcanoes: Still Alive After All These Years? space.com http://www.space.com/scienceastronomy/solarsystem/mars_volcano_011113.html. European Space Agency's Mars Express Mission homepage and image browser. Hartmann, W. K. and Neukum, G. (2001) Cratering chronology and the evolution of Mars. Space Science Reviews, v. 96, p. 165-194. Hauck, S. A. and Phillips, R. J. (2002) Thermal and crustal evolution of Mars. Journal of Geophysical Research, v. 107, doi 10.1029/2001JE0011801. Introduction to Cratering Studies and the Crater Counting Technique from the Planetary Science Institute, Tucson, Arizona. Kiefer, W. S. (2003) Melting in the Martian mantle: Shergottite formation and implications for present-day mantle convection on Mars. Meteoritics and Planetary Science, v. 38, p. 1815-1832.

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PSRD: Recent Activity on Mars: Fire and Ice

http://www.psrd.hawaii.edu/Jan05/MarsRecently.html

Milkovich, S M. and Head III, J. W. (2003) Olympus Mons Fan Shaped Deposit Morphology: Evidence for Debris Glaciers. Sixth International Conference on Mars, 3149 (pdf file). Neukum, G., Jaumann, R., Hoffmann, H., Hauber, E., Head J. W., Basilevsky, A. T., Ivanov, B. A., Werner, S. C., van Gasselt, S., Murray, J. B., McCord, T., and the HRSC Co-investigator team. (2004) Recent and episodic volcanic and glacial activity on Mars revealed by the High Resolution Stereo Camera. Nature, v. 432, p. 971-979. Neukum, G., Ivanov, B. A., and Hartmann, W. K. (2001) Cratering records in the inner solar system in relation to the lunar reference system. Space Science Reviews, v. 96, p. 55-86. Nyquist, L. E. et al, (2001) Ages and geologic histories of Martian meteorites. Space Science Reviews, v. 96, p. 105-164. Mars Odyssey GRS results are about to be submitted for publication, but see Taylor, G. J. et al. (2003) Igneous and aqueous processes on Mars: Evidence from measurements of K and Th by the Mars Odyssey Gamma Ray Spectrometer. Sixth International Conference on Mars, 3207 (pdf file).

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