A review of meteorite evidence for the timing of magmatism and of surface or near-surface liquid water on Mars

JOURNAL OF GEOPHYSICAL RESEARCH, VOL. 110,, doi:10.1029/2005je002402, 2005 A review of meteorite evidence for the timing of magmatism and of surface or near-surface liquid water on Mars Lars Borg Institute of Meteoritics, University of New Mexico, Albuquerque, New Mexico, USA Michael J. Drake Lunar and Planetary Laboratory, University of Arizona, Tucson, Arizona, USA Received 13 January 2005; revised 3 March 2005; accepted 6 April 2005; published 27 September 2005. [1] There is widespread photogeological evidence for ubiquitous water flowing on the surface of Mars. However, the age of surface and near-surface water cannot be deduced with high precision from photogeology. While there is clear evidence for old and young fluvial features in the photogeologic record, the uncertainty in the absolute calibration of the Martian crater flux results in uncertainties of ±1.5 Gyr in the middle period of Martian geologic history. Aqueous alteration of primary igneous minerals produces secondary minerals in Martian meteorites. Here we use the ages of secondary alteration minerals in Martian meteorites to obtain absolute ages when liquid water was at or near the surface of Mars. Aqueous alteration events in Martian meteorites occurred at 3929 ± 37 Ma (carbonates in ALH84001), 633 ± 23 Ma (iddingsite in nakhlites), and 0 170 Ma (salts in shergottites). Furthermore, these events appear to be of short duration, suggesting episodic rather than continuous aqueous alteration of the meteorites. The Martian meteorites appear to be contaminated by Martian surface Pb characterized by a 207 Pb/ 206 Pb ratio near 1. Lead of this composition could be produced by water-based alteration on the Martian surface. The high 129 Xe/ 132 Xe ratio in the Martian atmosphere compared to Martian meteorites indicates fractionation of I from Xe within 100 Myr after nucleosynthesis of 129 I. Such fractionation is difficult to achieve through magmatic processes. However, water very efficiently fractionates I from Xe, raising the intriguing possibility that Mars had a liquid water ocean within its first 100 Myr.1. Citation: Borg, L., and M. J. Drake (2005), A review of meteorite evidence for the timing of magmatism and of surface or near-surface liquid water on Mars, J. Geophys. Res., 110,, doi:10.1029/2005je002402. 1. Introduction [2] Mars is one of two terrestrial planets that have evidence for the presence of liquid water on its surface [e.g., Carr, 1996]. For example, recent investigations by the Mars Exploration Rover at the Opportunity landing site have revealed layered deposits that are interpreted to be sedimentary rocks that were produced by weathering of basalts [Squyres et al., 2004; Herkenhoff et al., 2004]. Outcrops composed of sulfate minerals that are thought to be of aqueous origin are also observed [Christensen et al., 2004; Klingelhöfer et al., 2004; Rieder et al., 2004]. The presence of water on Mars is also clearly manifested by fluvial features covering portions of the surface of Mars that are in many ways similar to those observed on Earth. The most common Martian fluvial features include incised valleys of variable length, often with multiple branched tributaries (valley networks), and long, wide channels emanating from broken and dissected chaotic terrain (outflow channels). The valley networks and outflow channels Copyright 2005 by the American Geophysical Union. 0148-0227/05/2005JE002402 are primarily concentrated in the oldest Noachian age and intermediate Hesperian age terrains, respectively. In contrast, the youngest Amazonian age terrain has significantly fewer of these fluvial features [e.g., Carr, 1996], leading to the hypothesis that Mars has become dryer through time. [3] The Martian meteorites also provide constraints on the history of water on Mars. Some of these samples record aqueous geochemical fractionation events that occurred on the surface and in the atmosphere and they preserve secondary alteration products produced by interaction of surface rocks with water. By analyzing both primary igneous minerals and secondary alteration products present in the Martian meteorites using a variety of isotopic systems, the timing of these water-based geochemical processes can be precisely determined. Despite the fact that these studies are limited to only a few individual samples, they provide the only mechanism to absolutely date Martian geologic processes associated with water. Therefore these ages are of fundamental importance in understanding the broader role of water on Mars. [4] Although both photogeologic and meteorite studies indicate a clear role of surface water in the Martian geologic record, the details of these studies have not penetrated both 1of10

scientific communities. This manuscript reviews the temporal constraints placed on these geochemical processes from radiogenic isotope studies of the Martian meteorites and attempts to correlate conclusions based on the radiogenic isotopic record of the meteorites with those from photogeology studies. In order to provide a geologic background in which to interpret the timing of water-based processes preserved in the meteorites, we will first review the absolute ages of geologic events on Mars inferred from the meteorites, followed by age constraints on the timing of aqueous geochemical processes on Mars as deduced from internal isochrons (e.g., Rb-Sr and Pb-Pb) of the meteorites. Finally we review the somewhat more indirect evidence for very early widespread water on Mars. This evidence comes from U-Pb and Xe isotopic measurements of Martian meteorites and from Xe isotopic measurements of the Martian atmosphere by the Viking spacecraft. 2. Characteristics of the Martian Meteorites [5] The goal of this review is to summarize the temporal constraints that can be placed on the history of water on Mars through chronologic studies of the Martian meteorites (shergottites, nakhlites, Chassigny, and ALH 84001). Because alteration of igneous rocks often results in the formation of secondary alteration products, and because it is the dating of these minerals that provides many of the temporal constraints discussed below, it is important to briefly review some of the pertinent petrologic features of this suite of samples. The meteorites represent a range of rock types with variable mineral modes and compositions. The shergottites are the most abundant class of Martian meteorite. These are basalts and basaltic cumulates containing some combination of pyroxene + plagioclase ± olivine [McSween, 1994] with ages less than 575 Ma [Borg et al., 2001]. They also contain various trace phases such as chromite, phosphate, ilmenite, pyrrhotite, silica, kaersutite, sulfides, carbonate, and sulfate, as well as variable amounts of impact derived melt. The nakhlites, the next most abundant class, are clinopyroxenites containing clinopyroxene, olivine, and mesostasis with minor to trace amounts of plagioclase, K-feldspar, phosphates, impact glass, magnetite, ilmenite, sulfates, carbonates, and the hydrous alteration material called iddingsite. Chassigny is a dunite, composed primarily of olivine with minor amounts of pyroxene, plagioclase, K-feldspar, biotite, kaersutite amphibole, baddeleyite, apatite, and sulfides. The nakhlites and Chassigny have ages of 1330 Ma [Nyquist et al., 2001]. The final meteorite type is represented by the orthopyroxenite ALH84001. In addition to orthopyroxene, this meteorite contains small amounts of olivine, plagioclase, K-feldspar, chromite, phosphates, silica, sulfides, magnetite, and carbonates, and has an age of 4.5 Ga [Nyquist et al., 1995]. From this description it is apparent that the Martian meteorites have a wide range of ages and contain minerals including the carbonates, iddingsite, sulfates, and clays that are interpreted to form as a result of aqueous alteration on Mars. 3. Timing of Geologic Events on Mars [6] In order to place aqueous geochemical processes into the context of the Martian geologic record, it is necessary to discuss the existing chronologic constraints on geologic processes derived from the meteorite suite. These include the timing of planetary differentiation (i.e., segregation of core, mantle, and crust) as well as the timing of magmatic activity. Figure 1 is a summary of ages determined for significant Martian geologic events. From Figure 1 it is apparent that the events recorded in the Martian meteorite suite occurred over discrete and limited time intervals. This ultimately reflects the fact that the meteorites are not a complete representation of surface rocks on Mars. Nevertheless, several age constraints on early geologic processes can be gleaned from the samples. [7] The age of core formation is based on the 182 Hf- 182 W (t 1/2 = 9 Ma) isotopic systematics observed in the meteorites. Small excesses of 182 W over other W isotopes indicate that W was fractionated from Hf at a time when 182 Hf was still alive. This fractionation event is interpreted as core formation because this process can strongly separate these elements. The age of core formation is estimated to be 4556 ± 1 Ma (Figure 2) assuming the Martian mantle has a Hf/W ratio of 5.1 [Foley et al., 2004, 2005]. The age of silicate differentiation is based on both the 146 Sm- 142 Nd (t 1/2 = 103 Ma) and 147 Sm- 143 Nd (t 1/2 = 106 Ga) isotopic systems. Enrichments and depletions of 142 Nd and 143 Nd observed in the meteorites indicate that Sm was fractionated from Nd at 4526 +21 19 Ma (Figure 2, revised from Borg et al. [2003a] using data from Foley et al. [2005]). This fractionation is interpreted to be the time that the mantle source regions of the basaltic Martian meteorites obtained their present-day compositions and hence represents the age of silicate differentiation [Harper et al., 1995; Borg et al., 1997, 2003a]. The observation that ancient ages of planetary differentiation are recorded in relatively young rocks indicates that the broad compositional and mineralogic features of major Martian reservoirs (e.g., core, mantle, and crust) were not only produced very early, but have been preserved throughout the history of the planet. Thus large-scale mixing processes associated with tectonics or giant impacts capable of producing magma oceans or lakes could not have occurred after 4.52 Ga [Borg et al., 1997; Foley et al., 2005]. [8] Other temporal constraints placed on Martian geologic processes are derived from the crystallization ages of the meteorites. The oldest age of 4500 ± 130 Ma is determined for ALH84001 [Nyquist et al., 1995]. The nakhlites and Chassigny have ages that are within analytical uncertainty of one another and yield a weighted average Sm-Nd age of 1327 ± 39 Ma (Figure 1). These are the next youngest Martian samples. The absence of samples with ages between 1.3 and 4.5 Ga is puzzling given the relatively large amount of Martian volcanic surface estimated to be of this age [Tanaka et al., 1992; Hartmann and Neukum, 2001; Neukum et al., 2001]. It could reflect the fact that the meteorites are only sampling a limited number of sites [Nyquist et al., 1998] or could reflect the difficulty of launching shock fragmented and/or altered material of this age off the Martian surface [Hartmann and Neukum, 2001; Hartmann, 2004]. From Figure 1, it is apparent that the shergottites define discrete age groupings. In addition to individual samples with ages of 575 and 474 Ma, there are three samples with 2of10

ages of 330 Ma and ten with ages of 175 Ma. The weighted average of the average Rb-Sr and Sm-Nd ages determined for individual samples from the youngest shergottite group yields an age of 174 ± 2 Ma. The small error associated with this calculation indicates that the ages of these samples cannot be distinguished and provides a minimum age spread for the time that these magmas were crystallized. It also suggests that these rocks are derived from a common igneous province on Mars. [9] The large number of young samples indicates that Mars is likely to be producing a significant volume igneous rock relatively recently, with magmatism possibly continuing today. The frequency of magma production can be estimated from the age range represented by the ten 175 Ma shergottites. In order to complete this calculation the number of individual magmas represented by the meteorites needs to be constrained and the age range of crystallization for the magmas must be defined. The range of initial Sr and Nd isotopic compositions of the handful of analyzed shergottites is larger than that observed for all terrestrial rocks [Borg et al., 2003a], and therefore provides the clearest mechanism to distinguish individual parental magmas (i.e., discrete igneous events). In fact, initial Sr and Nd isotopic compositions of all ten of these shergottites (including EET79001A and B) are unique, indicating that they are derived from ten different parental magmas (Figure 3). The duration of magmatism represented by these samples is more difficult to estimate. On one hand, this duration could be represented by the uncertainty Figure 1. Absolute ages of dated Martian events. The top panel shows ages of secondary alteration events (see text). The bottom panel shows crystallization and planetary differentiation ages. Shergotty, EET79001A, and Y793605 are Rb-Sr ages, whereas NWA1195, NWA1068, DaG 476, Dhofar 019, the nakhlites, and ALH84001 are Sm-Nd ages only. Other shergottite ages are weighted average of Rb-Sr and Sm-Nd ages. Data sources: Shergotty [Shih et al., 1982]; Zagami [Borg et al., 2003b, 2005]; LA1 [Nyquist et al., 2000]; NWA856 [Brandon et al., 2004]; EET79001A [Wooden et al., 1982]; Y793605 [Morikawa et al., 2001]; EET79001B [Nyquist et al., 2001]; ALH77005 and LEW88516 [Borg et al., 2002]; NWA1068 [Shih et al., 2003]; Y980459 [Shih et al., 2004]; QUE94201 [Borg et al., 1997]; NWA1195 [Symes et al., 2005]; DaG 476 [Borg et al., 2003a]; Dhofar 019 [Borg et al., 2001]; Nakhla [Nakamura et al., 1982]; NWA998 [Carlson and Irving, 2004]; Y000593 [Shih et al., 2002]; Lafayette [Shih et al., 1998]; Chassigny [Jagoutz, 1996]; Governador Valadares [Shih et al., 1999]; and ALH84001 [Nyquist et al., 1995]. Note that the age of core segregation is relative to solar system formation age of 4567 Ma [Amelin et al., 2002], and silicate differentiation age is relative to the age of the Angrite LEW86010 of 4558 Ma determined by Nyquist et al. [1994]. Core segregation and silicate differentiation represent absolute ages but must be linked to solar system formation and Angrite crystallization ages because the initial 182 Hf/ 183 W and 146 Sm/ 144 Nd used in these short-lived chronometers have been determined on materials with ages of 4567 and 4558 Ma, respectively. Figure 2. Two stage Nd isotope evolution model after Borg et al. [2003a]. Model isochron defines an age of silicate differentiation of 4526 +21 19 Ma relative to LEW86010 age of 4558 Ma, assuming initial 146 Sm/ 144 Nd ratio of 0.0076, and using Sm-Nd data from references cited in Figure 1 and Harper et al. [1995] and Foley et al. [2004, 2005]. 3of10

mechanisms that could account for the presence of significant amounts of secondary alteration products in the Martian meteorite suite. Figure 3. Initial 87 Sr/ 86 Sr ratio versus age for 175 Ma shergottites. Initial Sr and Nd isotopic compositions of individual meteorites differ outside analytical uncertainty indicating that 10 different magmatic events are represented by these samples. The range of eruptive ages is between 4 Myr (the weighted average of shergottite ages) and 42 Myr (the range of observed ages ± analytical uncertainty). These ranges correspond to magmatic rates of 1 event every 0.4 Myr to 4 Myr, respectively. determined for the weighted average of the shergottite ages (i.e., 4 Myr). On the other hand, the duration of magmatic activity for the 175 Ma shergottites could be represented by the spread of ages determined for the samples themselves ± the analytical uncertainty (i.e., 42 Myr). If the age range is 4 Myr, there must be at least one discrete magmatic event occurring on Mars every 0.4 Myr. Alternatively, if the age range is 42 Myr, then at least one magmatic event is occurring on Mars about every 4 Myr (Figure 3). Keep in mind that these represent maximum time intervals because they assume that every magmatic event occurring on Mars is represented by dated samples currently in hand. In all likelihood, most Martian magmatic events are not represented in the current Martian meteorite suite and magmatic events have occurred at a frequency significantly greater than every 4 Myr. As a result, Mars appears to have been extremely active, producing numerous magmas around 175 Ma. If magmatism is not episodic, this observation suggests that there should be numerous young intrusive and extrusive rocks near the surface of Mars. [10] A high frequency of recent magmatic events is potentially central for understanding mechanisms to produce secondary alteration products observed in the meteorites. This stems from the fact that these events could provide the heat necessary to mobilize fluids on the surface, or in the shallow subsurface, producing secondary mineralization throughout Martian history. In addition, large impacts can produce hydrothermal systems that could also result in the production of the secondary minerals observed in the meteorites [Newsom, 1980; Newsom et al., 2001]. In either case, there appears to be at least two 4. Age Constraints on Secondary Alteration [11] There is abundant evidence for secondary alteration of the Martian meteorites. In general, the formation of these secondary minerals is thought to be the result of evaporation of brines at temperatures between 25 to 150 C, which resulted in relatively simple mineral assemblages [Bridges et al., 2001]. The small amount of silicate alteration indicates that the alteration events were relatively short-lived and likely occurred in single episodes [Bridges et al., 2001], a necessary prerequisite for application of radiometric dating techniques. Below we review some of the most salient features of the alteration products in the context of radiometric dating studies. For a detailed description of secondary mineral assemblages, see Bridges et al. [2001]. 4.1. Carbonates in ALH84001 [12] Several of the Martian meteorites contain minerals that are interpreted to be produced by secondary alteration in the presence of water. The most notable secondary alteration products are the carbonates in ALH84001, which make up about 1 percent of the rock [Romanek et al., 1994, 1995]. The carbonates appear to have precipitated on surfaces and in fractures. In many places the carbonates are zoned from Ca- and Fe-rich in the cores to Mg-rich at the rims and contain both magnetite and Fe-sulfides [e.g., Harvey and McSween, 1996]. There has been wide speculation about the mechanisms by which the carbonates were produced, because the compositions of many of the carbonates fall in a miscibility gap between Ca- and Fe-Mg carbonates. One particularly insightful set of experiments was able to reproduce most of the mineralogic and compositional features observed in the carbonates [Golden et al., 2001]. These authors demonstrated that Ca-Fe-Mg zoned carbonates, as well as the associated magnetite and Fe-sulfides, can be produced by rapid precipitation from an aqueous solution at low temperature followed by mild heating to 470 C. These experiments offer strong evidence that the carbonates in ALH84001 were produced by relatively low temperature hydrous alteration processes. [13] The age of carbonate formation has been determined by Borg et al. [1999] using the Rb-Sr and Pb-Pb dating methods on a set of carbonate leachates. The Rb-Sr system yielded an age of 3900 ± 40 Ma, whereas the Pb-Pb system yielded an age of 4040 ± 100 Ma (Figure 4). The concordance of these ages suggests that the isochrons represent the formation age of the carbonates and are not the result of mixing processes associated with the leaching of the sample. This date strongly suggests that aqueous near surface alteration processes were occurring on Mars at 3929 ± 37 Ma (the weighted average of the Rb-Sr and Pb-Pb ages). Interestingly, the observation that the isotopic systematics of the carbonates appears to have remained undisturbed since this time indicates that the meteorite was not subjected to later alteration processes. By inference it seems water was present at the ALH84001 site for a limited duration early in the history of the planet but almost 600 Myr after 4of10

solubilities in water and may have been acquired by trapping of such water into the structure. If either of these interpretations is correct, ages determined on iddingsite formation potentially date periods of water-based alteration occurring in the nakhlites. [15] There have been two approaches used to date the time of iddingsite formation in the nakhlites (Figure 5). The first is based on the Rb-Sr isotopic systematics of iddingsite-rich mineral fractions and leachates. Tie lines constructed between these fractions in Lafayette have a slope corresponding to an age of 679 ± 66 Ma [Shih et al., 1998]. Likewise ages of 650 ± 80 Ma and 614 ± 29 Ma have been defined for iddingsite-rich olivine residue and olivine leachate fractions from the nakhlite Y000593 by Shih et al. [2002]. The second approach to date iddingsite formation is based on the K-Ar system. Swindle et al. [2000] determined an age of 670 ± 91 Ma on a single iddingsite-rich sample from Lafayette. Although other K-Ar ages determined on similar samples were younger than 670 Ma, Swindle et al. [2000] argued that Ar had been lost from these samples by thermal metamorphism and that the most likely age of alteration was 670 Ma. The fact that three Rb-Sr ages and one K-Ar age determined on iddingsite-rich fractions from two different nakhlites are concordant strongly suggests that secondary alteration occurred around 650 Ma. If these ages represent a single alteration event occurring in the nakhlites on Mars, then the best age estimate for this event is the weighted average of these ages (633 ± 23 Ma; Figure 5). Figure 4. Rb-Sr and Pb-Pb isochrons for carbonate leachates from ALH84001 calculated from data by Borg et al. [1999]. Ages are slightly different from those reported by Borg et al. [1999] because they were calculated using the Isoplot program [Ludwig, 2001]. Weighted average of Rb- Sr and Pb-Pb ages is 3929 ± 37 Ma and is most representative of the age of secondary alteration in this meteorite. crystallization of ALH84001, and that the site has not been subjected to subsequent alteration. 4.2. Iddingsite in the Nakhlites [14] All of the nakhlites contain the hydrous alteration material that has been termed iddingsite. Iddingsite is particularly abundant in Lafayette where it has been documented to replace olivine, pyroxene, and high silica glass with smectite clays, magnetite/maghemite, and ferrihydrite [Bunch and Reid, 1975; Boctor et al., 1976; Treiman et al., 1993; Treiman and Lindstrom, 1997]. Veins of iddingsite are cross cut by the fusion crust in Lafayette, leading to the conclusion that iddingsite in the nakhlites formed as a result of aqueous alteration occurring on Mars [Treiman et al., 1993]. Drake et al. [1994] have argued that aqueous alteration is required to produce both the elevated 129 Xe/ 132 Xe ratios in Nakhla and iddingsite. Arriving at a consistent conclusion using different logic, Bogard and Garrison [1998] argued that Ar/Xe and Kr/Xe ratios in ALH84001 and, probably, the nakhlites resemble their Figure 5. Ages of iddingsite-rich fractions from the nakhlites Lafayette and Y000593. Open circles are K-Ar ages reported by Swindle et al. [2000] for Lafayette, whereas filled circles and squares are Rb-Sr ages determined by Shih et al. [1998, 2002] for both meteorites. The age of iddingsite formation is best represented by the weighted average of the three Rb-Sr ages and oldest K-Ar age to be 633 ± 23 Ma. Younger K-Ar ages are interpreted to have been reset by thermal metamorphism. 5of10

[16] Like the carbonate alteration in ALH84001, iddingsite formation in the nakhlites appears to have occurred at a single point in time. This suggests that alteration in these samples is episodic, otherwise the isotopic systematics of alteration products would not yield ages. The more limited style of alteration on Mars in comparison to Earth most likely reflects a more limited supply of water on the Martian near surface. As a consequence, it seems that alteration in ALH84001 and the nakhlites is not associated with continual weathering occurring on the surface, but rather with remobilization of water associated with discrete geologic events. Thus alteration is most likely associated with impacts or igneous activity. In this regard, it is perhaps significant to note that the age of iddingsite formation of 633 ± 23 Ma is nearly concordant with the 575 ± 7 Ma crystallization age of the shergottite Dhofar 019 [Borg et al., 2001]. 4.3. Secondary Minerals in the Shergottites [17] Numerous secondary minerals have been identified in the shergottite EET79001 and are summarized by Stoker et al. [1993]. These phases include: Ca, Mg, and Fe-Mn carbonates, Ca and Mg sulfates, Mg phosphate, Na-K chloride, illite, and smectite. Many of these minerals contain small amounts of water. Gooding et al. [1988, 1990] have argued that these minerals are of pre-terrestrial origin on the basis of textural evidence that suggests they were present in the meteorite prior to formation of shock features. Furthermore, Gooding et al. [1990] demonstrate that the concentrations of volatile species measured in EET79001, as well as ALH77005 and Shergotty, are consistent with formation of the sulfates, carbonates, and possibly chloride compounds, by aqueous geochemical processes occurring on Mars. The presence of these secondary alteration products in meteorites with young ages indicates that aqueous alteration must be occurring on Mars relatively recently, i.e., within the last 175 Myr. [18] Unlike the alteration observed in ALH84001 and the nakhlites, the alteration observed in the shergottites is on a significantly smaller scale. This style of limited alteration might result from continual weathering processes associated with the seasonal atmospheric water cycle on Mars. If this is the case, it suggests that formation of alteration products on Mars occurs as a result of discrete impact and/or magmatic events and global climatic processes. 5. Evidence for Surface Water From Isotopic Systematics of Martian Surface Material and Atmosphere [19] In addition to the formation of secondary alteration products in the Martian meteorites, fractionation of parent and daughter isotopes in Martian soils and atmosphere also suggest the involvement of liquid water in geochemical processes occurring on the surface of Mars. The basis for these arguments are the U-Pb isotopic systematics of the Martian meteorites and the I-Xe isotopic systematics of the Martian atmosphere. 5.1. U-Pb Isotopic Systematics of the Martian Meteorites [20] Uranium-lead isotopic analyses have been completed on several Martian meteorites [e.g., Chen and Wasserburg, 1986; Misawa et al., 1997; Borg et al., 2003b, 2005]. Defining the crystallization ages of the meteorites from these studies has proved a complicated task, because the meteorites contain a contaminant that has significantly altered their Pb isotopic systematics. Borg et al. [2003b, 2005] were able to obtain a precise 238 U- 206 Pb age for Zagami that was concordant with Rb-Sr and Sm-Nd ages determined on the same mineral fractions. The 235 U- 207 Pb age that they obtained, however, is disturbed leading to the suggestion that the contaminant observed in Zagami, and the other Martian meteorites, was characterized by an elevated 207 Pb/ 206 Pb ratio of at least 1.0 (modern terrestrial Pb has a 207 Pb/ 206 Pb ratio of 0.77 to 0.87). They also noted that the contaminant was most prevalent in impact glass-rich mineral fractions. Similar impact glass from EET79001 contains a significant excess of sulfur that is argued to result from the addition of soil to the meteorite during weathering-impact processes occurring on the Martian surface [Rao et al., 1999]. This is supported by mass balance calculations that indicate that up to 10% of the impact melt glass in EET79001 could be composed of soil with compositions matching those measured by the Viking and Pathfinder missions. It is therefore probable that the sulfur-rich soil observed on Mars is characterized by an elevated 207 Pb/ 206 Pb ratio of a least 1. [21] The production of material with a 207 Pb/ 206 Pb ratio of 1 is difficult to do by igneous processes because it requires significant fractionation of U from Pb and these elements behave similarly in most igneous systems. On the other hand, low temperature water-based alteration on Earth can produce large fractionations of U from Pb as a result of the greater solubility of U relative to Pb in these systems [e.g., Asmerom and Jacobsen, 1993]. Furthermore, chemical precipitates formed from aqueous solutions may also have low U/Pb ratios and thus fractionate these elements from one another. The soils could therefore represent material from which U was removed, and Pb concentrated, by water. In order to produce the observed 207 Pb/ 206 Pb ratio the fractionation of U from Pb must occur prior to 2 Ga. The more efficient the fractionation, the more recent the fractionation event. Therefore, if this scenario is correct, it suggests that the soils gained their geochemical characteristics prior to 2 Ga. This is consistent with photographic evidence that suggest that the Martian surface was significantly wetter in the past [e.g., Carr, 1996; Hartmann, 2003]. We therefore speculate that the chemical characteristics of the soil were set in the Noachian when water was most abundant on the Martian surface. 5.2. I/Xe Evidence for Liquid Water by 4.45 Gyr Ago [22] An intriguing observation is that Mars retains reservoirs reflecting different 129 Xe/ 132 Xe ratios. The Martian atmosphere has a high 129 Xe/ 132 Xe ratio compared to Martian meteorites, quite the opposite to the relationship of mid-ocean ridge basalts and the atmosphere on Earth (Table 1). The isotope 129 Xe is produced through the decay of now extinct 129 I(t 1/2 = 16 Ma). After 5 7 half lives, radiogenic daughter 129 Xe from the decay of 129 I cannot be detected. Hence any variation in 129 Xe/ 132 Xe ratios must be produced within approximately the first 100 million years after nucleosynthesis of 129 I. 6of10

Table 1. Ratios of 129 Xe to 132 Xe in Mars and Earth Reservoirs Mars Earth Atmospheres 2.4 0.985 Basalts 0.988 OIB a 1.0 1.5 1.00 1.14 MORB a a OIB refers to ocean island basalts; MORB refers to mid-oceans ridge basalts. [23] A 129 Xe/ 132 +2 Xe ratio of 2.5 1 for the Martian atmosphere was first determined from the Viking aeroshell mass spectrometer data [Owen et al., 1977] and has since been more accurately determined in gas-rich glassy inclusions from EET79001 [Swindle et al., 1986] to be about 2.4. The lowest 129 Xe/ 132 Xe ratio of 1.029 ± 0.019 has been measured in Chassigny [Ott and Begemann, 1985]. The true 129 Xe/ 132 Xe ratio for the Martian mantle is unknown. However, contamination of Martian meteorites with Martian atmosphere could only raise the measured ratio, so the Martian mantle ratio is unlikely to be higher than the Chassigny ratio [Musselwhite et al., 1991]. [24] There are two possible ways to fractionate I from Xe in the first 100 Myr of Martian history: magmatic processes and aqueous processes. In fact, it is conceivable that both processes played a role, although water is far more effective in fractionating I from Xe than mantle melting. 5.2.1. Magma Ocean [25] Musselwhite and Drake [2000] examined two bracketing magmatic scenarios that might account for the development of reservoirs with distinct 129 Xe/ 132 Xe ratios on Mars. Although both I and Xe are incompatible elements, the solubility of I in magmas is higher than the solubility of Xe. A high 129 Xe/ 132 Xe ratio in the atmosphere requires that the I/Xe ratio of the Martian mantle be raised because the mantle likely outgassed to form the atmosphere. In one scenario, an early magma ocean is outgassed with the attendant atmosphere having a low I/Xe ratio, because I is preferentially retained in the magma compared to Xe. The atmosphere was then stripped, perhaps by hydrodynamic escape and/or heavy bombardment, followed by secondary outgassing of a high I/Xe ratio magma ocean or solid mantle. Obviously this atmospheric stripping would have to occur while 129 I was still available, i.e., in less than about 100 Myr following nucleosynthesis of 129 I. Greater fractionation would occur with earlier outgassing, but the time for atmospheric stripping would be correspondingly shortened. In the second scenario, atmospheric stripping was simultaneous with primary magma ocean outgassing. Musselwhite and Drake [2000] showed that magmatic processes could account for the different Martian 129 Xe/ 132 Xe reservoirs, but only with exquisite coordinated timing of magma ocean and atmospheric stripping events. The details and variables in the analysis are complex and interrelated, and the reader is referred to Musselwhite and Drake [2000] for a full discussion. 5.2.2. Water Ocean [26] More promising, however, is the role of water. Iodine is highly soluble in water, whereas Xe is insoluble. This fractionation is extreme because I is a factor 10 11 more soluble in water than Xe [Musselwhite et al., 1991]. Outgas I and Xe into a liquid water ocean and I dissolves while Xe bubbles through to the atmosphere. Musselwhite et al. [1991] have presented a model to account for the Martian Xe data that relies on these very different solubilities of I and Xe in water to produce very efficient I/Xe fractionation during outgassing. Accretion ceases while 129 I still exists, the magma ocean solidifies, at least at the surface, and liquid H 2 O becomes stable. Iodine and Xe are both outgassed from the Martian mantle. Iodine dissolves in the ocean and is incorporated hydrothermally into the crust, whereas Xe remains in the atmosphere. Hydrodynamic escape and/or early heavy bombardment erodes the atmosphere, leaving most liquid water standing, so that I dissolved in the water is retained, but the xenon in the atmosphere is removed. In the hydrothermally altered crust 129 I decays to 129 Xe, which is released over geologic time, accounting for the high 129 Xe/ 132 Xe ratio (and high 40 Ar/ 36 Ar ratio) in the atmosphere compared to Martian meteorites. The strength of this hypothesis is that the 10 11 difference in solubility allows even widespread lakes or seas, rather than planetary-scale oceans, to provide the necessary fractionation of I from Xe. [27] The discussion above indicates that magmatic processes cannot be ruled out as the mechanism for different 129 Xe/ 132 Xe reservoirs on Mars. However, the extreme fractionation of I from Xe in liquid water compared to the modest fractionation in magmatic processes leads to the possibility that liquid water was stable at the Martian surface immediately following accretion and solidification of at least the outer part of the magma ocean. If accretion ended no later than 4.52 Ga [Borg et al., 2003a] water could have been present soon after. 6. Geologic History of Mars Inferred From the Meteorites and Photogeology 6.1. Ages of Martian Meteorites and Martian Surfaces [28] The young crystallization ages of the meteorites suggest that there is a significant amount of recent igneous activity on Mars. This conclusion is also derived from studies of surface ages based on crater counts [e.g., Hartmann, 1973]. In fact, lava flows less than 600 Ma (the age of the oldest shergottite) have been identified in several locations on Mars including the Elysium area [McEwen et al., 1999; Hartmann, 1999; Hartmann and Berman, 2000], the caldera of Arsia Mons [Hartmann et al., 1999], Amazonis Planitia [Hartmann, 1999], and Cerberus Fossae [Berman and Hartmann, 2002; Burr et al., 2002]. Although the abundance of Martian meteorites with young ages is consistent with young surface ages determined from photogeologic studies, the proportion of young meteorites is not. This stems from the fact that all but one of the twenty-two dated Martian meteorites have crystallization ages 1.3 Ga, yet only a small percentage of the Martian surface is this young [Tanaka et al., 1992; Nyquist et al., 1998]. Furthermore, the absence of meteorites with ages of intermediate age, if taken at face value, suggests that igneous rocks with ages between 1.5 and 4.0 Ga are either not produced on Mars or were not preserved on Mars. Thus the ages of the Martian meteorites paint a picture of an extremely young planet that is exceedingly active geologically. On the other hand, the large proportion of highly cratered terrain suggests a significantly less active planet. 7of10

[29] The overrepresentation of young meteorites has led to the suggestion that several meteorites are derived from single ejection sites [e.g., Nyquist et al., 1998]. The frequency of individual exposure ages determined on the Martian meteorites has been used as a proxy for the number of ejection events of the meteorites off of the surface of Mars. Eugster et al. [2002] noted that there is a range of ejection ages from 0.7 to 20 Ma. Using a combination of ejection ages, crystallization ages, and petrologic characteristics of the meteorites it has been concluded that there are roughly 4 or 5 discrete Martian sites represented in the meteorite suite [e.g., Nyquist et al., 1998, 2001; Eugster et al., 2002]. Although this model decreases the apparent overrepresentation of young meteorites, all but one of these sites must be on terrain with an age less than 1.3 Ga (i.e., 75 80% of the impacted terrain must be 1.3 Ga). The large proportion of young meteorites is perhaps the greatest discord between the geologic history of Mars inferred from the meteorites and that inferred from photogeologic studies and suggests that either (1) young rocks are preferentially ejected from Mars [Hartmann, 2004], (2) the Martian surface is significantly younger than we think it is [Nyquist et al., 1998], (3) there are less than 4 5 ejection events, or (4) some combination of all three. 6.2. Recent Surface Water on Mars and Radiometric Ages of Secondary Alteration [30] Age dating of secondary alteration in the Martian meteorites demonstrates that short duration water-based alteration events have occurred throughout Martian history. The presence of ancient alteration products in ALH84001 is consistent with the abundance of valley networks on ancient Noachian terrain that indicate that early Mars was a relatively wet place. Recent alteration observed in the nakhlites and shergottites is also consistent with photogeologic evidence for extremely recent outflow channels in some areas of Mars. For example, some lava flows in Marte Vallis that have been dated using crater counting to less than 200 Ma [Berman and Hartmann, 2002] have also been cut by channels [Burr et al., 2002]. These observations confirm that even the youngest igneous rocks on Mars have undergone water-based erosion, and hence potentially aqueous alteration. [31] It is interesting to note that the amount of secondary mineralization that is present in the meteorites appears to decrease toward the present-day. The 4500 Ma ALH84001 is made up of 1% carbonates, the 1330 Ma nakhlites have obvious iddingsite alteration products that make up trace amounts of the meteorites, whereas some of the shergottites have vanishingly small amounts of various salts and clays. Furthermore, the alteration of each meteorite appears to have occurred over a discrete period of time. Therefore, although not definitive by any means, this correlation is certainly consistent with the hypothesis that Mars has become dryer through time. [32] The Xe isotopic systematics of the meteorites can be interpreted to imply the presence of an ancient Martian ocean. This is consistent with several photogeologic studies that conclude that oceans were once present on Mars [e.g., Baker et al., 1991; Parker et al., 1993; Edgett and Parker, 1997; Fairén et al., 2003]. It has been suggested that fossil shorelines may be present in Vastitas Borealis region and represent a primordial ocean that could have existed at a time when 129 I was still alive [Clifford and Parker, 2002]. This suggestion is also consistent with the observation that rampart craters and other features characteristic of large amounts of ground ice have been observed in this region [Hartmann, 2003]. 7. Conclusions [33] There is definitive evidence in the Martian meteorite suite for igneous events at 4500 ± 130 Ma, 1327 ± 39 Ma, 575 ± 7 Ma, 474 ± 11 Ma, 332 ± 9 Ma, and 174 ± 2 Ma. Aqueous alteration events occurred at 3929 ± 37 Ma (carbonates in ALH84001), 633 ± 23 Ma (iddingsite in nakhlites), and 0 170 Ma (salts in shergottites). The discrete time intervals of aqueous alteration recorded in these classes of Martian meteorites and their separation in time from the igneous events that produced the rocks argues for episodic, rather than continuous, presence of surface or near surface water. Martian meteorites cannot be considered a representative sampling of the Martian surface. Mars must have had short-lived pulses of aqueous alteration events more frequently than recorded in the meteorite suite. These events may have been associated with magmatic activity which may have mobilized water ice. 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