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The Role of Halogens During Fluid and Magmatic Processes on Mars

  • Elizabeth B. Rampe
  • Julia A. Cartwright
  • Francis M. McCubbin
  • Mikki M. Osterloo
Chapter
Part of the Springer Geochemistry book series (SPRIGEO)

Abstract

The geochemistry of halogens on Mars gives insight into the composition of the martian mantle, igneous evolution of the martian crust, aqueous processes on the martian surface, and the overall habitability of the planet. Halogen abundances have been measured from martian meteorites, in situ by landers and rovers, and from orbital missions around Mars. The bulk rock abundances of halogens have been determined for many martian meteorite samples including all petrological five types (nakhlites, chassignites, shergottites, orthopyroxenites, and regolith breccias). Measurements of basaltic martian meteorites (i.e., shergottites and regolith breccias) provide important insights into halogen abundances in mantle and crustal reservoirs. Fluorine, Cl, Br, and I have been detected in silicate, phosphate, sulfate, oxide, and halide group minerals in martian meteorites. These halogen-bearing minerals are found in melt inclusions, as secondary hydrothermal or aqueous alteration products, or in the interstices between cumulus igneous silicates. Measurements from meteorites and from martian missions indicate that Cl is the most abundant halogen on and in Mars. Measurements from landed missions suggest that Cl is commonly present in oxychlorine compounds (e.g., perchlorate and chlorate salts), whereas measurements from orbit have identified both oxychlorine minerals and halite. Halite is constrained to local depressions in the ancient southern highlands, suggesting precipitation from the evaporation of water in closed basins at ~3.5–4 Ga. The presence of oxychlorine minerals on the martian surface has important implications for the habitability of present day Mars because oxychlorine minerals may deliquesce to create seasonal deposits of liquid water. Furthermore, oxychlorine compounds are considered both a resource and potential hazard to the eventual human exploration of Mars.

16.1 Introduction

Other than the Earth and the Moon, Mars may be the best-studied object in the Solar System. Geological and geochemical data have been collected by orbiters and landers, and a wealth of knowledge has been gained through the study of martian meteorites. Mars’ early history may have mimicked that of the Earth, including significant volcanism and an active hydrological cycle. There is abundant geomorphological and mineralogical evidence for surface or near-surface water early in Mars’ history (e.g., Carr 1995; Ehlmann et al. 2011). Life as we know it on Earth requires liquid water, and the relative proximity of Mars to Earth makes it a high-priority target for answering the question: “Are we alone in the universe?”

Mars, however, differs in many ways from the Earth. There is little evidence for plate tectonics, and, as such, there may not have been a planet-wide process to recycle the crust and create evolved magmas. Regarding magnetism, Mars lost its dynamic magnetic field ~4 billion years ago (Ga) following the termination of an internal dynamo (e.g., Acuña et al. 1999). As a result, the atmosphere was gradually stripped by solar-wind-induced sputtering and photochemical escape (e.g., Jakosky et al. 1994), causing the loss of much of Mars’ surface water. The noble gas signature of the martian atmosphere hints at a significant loss early in its history, leading to elevations in certain noble gas isotope abundances (e.g., Jakosky et al. 1994; Jakosky and Jones 1997). Mars is also geochemically different from Earth because it accreted from a more volatile-rich portion of the protoplanetary disk. The silicate portion of Mars is more enriched in FeO, MnO, alkalis, and halogens and is depleted in siderophile and chalcophile elements (Longhi et al. 1992). Mars also has an iron-nickel-sulfur (Fe-Ni-S) core (Longhi et al. 1992), compared to Earth’s core, which is mostly dominated by Fe and Ni.

Many details surrounding Mars’ early history remain a mystery. Halogens, however, can give us an insight into martian mantle composition, igneous evolution, aqueous processes, and climatic history. In this chapter, we review both the abundance and mineralogy of halogens in martian meteorites and discuss the significance of these measurements for the martian interior and crust. We then review the detection of halogens on the surface from landed and orbital missions and consider the speciation of halogens and the implications for climate and habitability.

16.2 Halogens in Martian Meteorites

The halogens—fluorine (F), chlorine (Cl), bromine (Br), and iodine (I) make up group 7 of the periodic table and are considered to be moderately to highly volatile. Early studies of martian meteorites highlighted that Mars is a volatile-rich planet (Dreibus and Wanke 1985). The bulk rock abundances of halogens in martian meteorites have been determined for a number of samples across all five martian meteorite types (i.e., nakhlites, chassignites, shergottites, orthopyroxenites, and regolith breccias). Bulk rock and mineral separate/chemical aliquot abundances for halogens within martian meteorites are shown in Table 16.1, and plotted in Fig. 16.1. The overall martian halogen abundances indicate that, in general, Cl is more abundant than F, and both are more abundant than Br and I, which is in-keeping with the relative abundances of halogens in CI chondrites (Lodders 2003). As bulk rock compositions of cumulate rocks do not represent magmatic liquids, halogen bulk rock abundances of martian cumulate rocks, like the chassignites, nakhlites, and orthopyroxenite (Allan Hills 84001), provide little insight into the halogen abundances of the martian interior. However, basaltic rocks, like the shergottites, and highly mixed sedimentary rocks like the regolith breccia Northwest Africa 7034 and its pairings, could provide important insights into halogen abundances in mantle and crustal reservoirs.
Table 16.1

Summary of halogen data for martian meteorites

Meteorite

F (ppm)

Cl (ppm)

Br (ppb)

I (ppb)

Reference

Shergottites

ALHA 77005

22

14

69

1720

Dreibus and Wänke (1987)

DaG 476

840

Zipfel et al. (2000)

Dho 019

 

22

Williams et al. (2016)

EET 79001 Lith A

39

26

189

100

Dreibus and Wänke (1987)

EET 79001 Lith A

 

180

Williams et al. (2016)

EET 79001 Lith B

31

48

287

960

Dreibus and Wänke (1987)

EET 79001 Lith B

 

42

  

Williams et al. (2016)

EET 79001 Lith C

 

35

378

12

Dreibus and Wänke (1987)

LEW 88516

27

29

Lodders (1998)

Los Angeles

 

115

Williams et al. (2016)

NWA 1068

13

Bogard et al. (2010)

NWA 2975

 

24

Williams et al. (2016)

NWA 6234

59

Burgess et al. (2013)

QUE 94201

40

91

Dreibus et al. (1996)

RBT 04262-Ol/Pyx

14 ± 2

Cartwright et al. (2009)

RBT 04262-Maskelynite

47 ± 1

Cartwright et al. (2009)

RBT 04262-Accessory

175 ± 52

Cartwright et al. (2009)

Shergotty

41.6

108

890

36

Dreibus and Wänke (1987)

Shergotty

 

108

Williams et al. (2016)

SaU 005

56

143

Dreibus et al. (2000)

Tissint-Bulk

 

22

Williams et al. (2016)

Tissint-Glass

 

69

Williams et al. (2016)

Tissint-Igneous

 

30

Williams et al. (2016)

Y-980459

86

60.7

Dreibus et al. (2003)

Zagami

41

145

760

4

Dreibus and Wänke (1987)

Zagami

 

49

Williams et al. (2016)

Nakhlites

Lafayette

 

101 ± 5

590 ± 30

54 ± 3

Dreibus et al. (2006)

Lafayette Iddingsite

 

3600

9000 ± 3000

Treiman and Lindstrom (1997)

MIL 03346

 

248 ± 12

450 ± 23

1590 ± 80

Dreibus et al. (2006)

MIL 03346

 

163.7 ± 0.9

411.8 ± 2.2

15.1 ± 0.6

Cartwright et al. (2013)

Nakhla

 

563

3410

24

Dreibus and Wänke (1987)

“Nakhla E”

 

1890 ± 95

3460 ± 173

17 ± 1

Dreibus et al. (2006)

“Nakhla K”

 

1145 ± 57

4300 ± 215

10 ± 1

Dreibus et al. (2006)

“Nakhla G”

 

876 ± 44

8450 ± 423

26 ± 1

Dreibus et al. (2006)

Nakhla Vein

 

7700 ± 900

151000 ± 49000

Rao et al. (2005)

Nakhla–Bulk Crush

 

104.5 ± 4.9

952.7 ± 5.3

11.8 ± 0.5

Cartwright et al. (2013)

Nakhla–Olivine Crush

 

14.2 ± 0.3

557.9 ± 2.0

20.0 ± 0.4

Cartwright et al. (2013)

Nakhla–Pyroxene Crush

 

59.2 ± 4.5

722.2 ± 1.6

12.3 ± 0.3

Cartwright et al. (2013)

Nakhla–Mesostasis

 

75.7 ± 2.1

Cartwright et al. (2013)

Nakhla–Pyroxene

 

40.5 ± 1.1

Cartwright et al. (2013)

Nakhla–Olivine

 

17.8 ± 0.6

Cartwright et al. (2013)

Nakhla–Bulk1

 

56.4 ± 3.5

Cartwright et al. (2013)

Nakhla–Bulk2

 

45.0 ± 0.9

Cartwright et al. (2013)

Nakhla–Bulk Acid-Etched

 

5.7 ± 0.2

Cartwright et al. (2013)

Nakhla–Bulk Water-Etched

 

35.3 ± 0.9

Cartwright et al. (2013)

NWA 817

 

117

Williams et al. (2016)

NWA 998

 

127 ± 6

180 ± 9

281 ± 14

Dreibus et al. (2006)

NWA 998–Feldspar

 

55.4 ± 1.7

Cartwright et al. (2013)

NWA 998–Pyx1

 

11.1 ± 0.3

Cartwright et al. (2013)

NWA 998–Pyx2

 

15.4 ± 0.2

Cartwright et al. (2013)

NWA 998–Pyx3

 

48.2 ± 1.3

85.6 ± 1.0

66.1 ± 1.4

Cartwright et al. (2013)

NWA 998–Pyx-4

 

35.5 ± 4.1

47.5 ± 2.2

105.6 ± 4.9

Cartwright et al. (2013)

NWA 998–Olivine

 

62.8 ± 8.8

45.7 ± 1.9

9.5 ± 1.1

Cartwright et al. (2013)

NWA 5790

 

72

Williams et al. (2016)

Y-000593

 

101 ± 5

80 ± 4

378 ± 19

Dreibus et al. (2006)

Y-000749

 

73 ± 4

60 ± 3

682 ± 34

Dreibus et al. (2006)

Chassigny

14.7

34

97

10

Dreibus and Wänke (1987)

NWA 7034

 

2200

Williams et al. (2016)

Analytical uncertainties are shown, where reported

Fig. 16.1

A plot showing the halogen abundances for F (ppm), Cl (ppm), Br (ppb), and I (ppb) determined from analyses of martian meteorites, as bulk values, mineral separates, and other chemical aliquots, as shown in Table 16.1. For references, see Table 16.1

One important point to make at this stage is that halogens can occur within meteoritic samples as either a primary component that was incorporated into the sample on formation, or they can become incorporated later due to secondary processes such as aqueous alteration. Bulk analyses are thus useful in giving an idea of the overall abundances present, but mineral separate analyses allow us to better characterize the processes leading to halogen occurrence within materials. There is some additional complexity to using bulk rock halogen abundances as a proxy for interior values. Volatile elements, in particular Cl, are susceptible to degassing (e.g., Aiuppa et al. 2009), which can result in underestimates of halogen abundances in parental magmas and magmatic source regions. In particular, studies by McCubbin and colleagues (2013–2016) have suggested that H and Cl degassing in martian samples is highly likely due to the fact that: (1) the bulk rock abundances of H2O and Cl are much lower than expected from reconstructing the abundances of H and Cl from mineral-melt partitioning relationships; (2) measurements of F and Cl in some magmatic apatite grains demonstrate that F has been concentrated; and (3) the Cl/F ratios in the martian mantle calculated from bulk rock measurements are less than terrestrial Cl/F ratios (McCubbin et al. 2013, 2015, 2016; McCubbin and Jones 2015; Filliberto et al. 2016).

While the abundances of halogens within meteorites are important, so too are the actual minerals that host them. Martian meteorites contain a diverse array of mineral phases that have halogens as essential structural constituents, commonly substituting for OH groups in mineral structures. For the majority of minerals, halogens are present as F and/or Cl members, though Br and I can occur in trace abundances. These minerals include silicates, phosphates, sulfates, oxides, and halides (Bunch and Reid 1975; Floran et al. 1977; Johnson et al. 1991; McSween and Harvey 1993; Watson et al. 1994; McSween and Treiman 1998; Leshin 2000; Righter et al. 2002; Boctor et al. 2003; Sautter et al. 2006; McCubbin and Nekvasil 2008; Greenwood et al. 2008; Filiberto and Treiman 2009a; McCubbin et al. 2009, 2010, 2012, 2013, 2016; Gross et al. 2013; Filiberto et al. 2014, 2016; Muttik et al. 2014; Santos et al. 2015). Understanding the origin of these halogen-bearing minerals is important for understanding both the petrogenesis of the martian meteorites and the behavior of halogens in martian magmatic, hydrothermal, and aqueous systems. Furthermore, they are important for constraining the abundances of halogens in the martian mantle and crust. In the following text, we will review the halogen-bearing minerals that have been identified in each of the distinct martian meteorite types, focusing on the distribution of F, Cl, and OH in the monovalent anion sites within each of the minerals that accept halogen elements as essential structural constituents.

Given that both the bulk rock abundances of halogens and halogen mineralogy are of critical importance for understanding the halogen budget of the martian interior, we have summarized the bulk rock and mineral data for each of the five martian meteorite groups below. These data will be discussed later in the chapter in order to understand the roles that halogens play in geologic processes in the martian interior and to place constraints on the halogen abundances in the martian crust and mantle.

16.2.1 Shergottites

The shergottite meteorites are the most abundant group of materials from Mars and are generally basaltic in composition. At the time of writing, there are 154 separate stones that constitute the shergottite meteorite group, though many of these are paired (they are genetically related, and likely split during impact). They are basaltic rocks that display a range of compositional and textural features relating to martian melt compositions with varying degrees of partial accumulation and/or fractionation. The shergottites are commonly grouped based on their abundances of incompatible trace elements and compositions with respect to Sm-Nd and Rb-Sr isotopic systems into three distinct groups, including enriched, intermediate, and depleted (Borg and Draper 2003). In addition, some studies have suggested that they may be genetically related to substantial volcanic systems (Symes et al. 2008). However, the bulk rock abundances of halogens do not correlate strongly with these geochemical groupings (Table 16.1; Filiberto et al. 2016). Of the major minerals present in shergottites, apatite—Ca5(PO4)3(OH,F,Cl)—(Fig. 16.2) is the only phase that is suitably halogen-bearing (McCubbin et al. 2016). Previous work has reported on the occurrence of kaersutitic amphibole (a Ti-rich variety of calcic amphibole) in some melt inclusions within shergottites (Treiman 1986), though these reports have not been confirmed by analysis of halogen or OH abundances. The apatites present in shergottites span a wide range of compositions that encompass fluorapatite, chlorapatite, and hydroxyapatite (Fig. 16.3). The occurrence of these minerals is important because the mineral compositions, bulk-rock abundances of F, and mineral-melt partitioning relationships can be used to quantify pre-eruptive (pre-degassing) water and Cl abundances in martian magmas, which have been used, in part, to constrain the halogen and water abundances of the martian mantle in previous studies (e.g., Gross et al. 2013; McCubbin et al. 2016).
Fig. 16.2

Back scattered electron (BSE) and elemental X-ray maps of apatite–merrillite intergrowth in Zagami. a BSE image of Zagami. All phases present are identified, and the phase abbreviations are indicated as follows: Ap apatite, Cpx clinopyroxene, Me merrillite, Msk maskelynite, Opx orthopyroxene, Si silica, and Tmt titanomagnetite. b Mg X-ray map, c Cl X-ray map, d Si X-ray map, e P X-ray map, f Ca X-ray map, and g Fe X-ray map

Fig. 16.3

Ternary plots of halogen-site occupancy (expressed as mol% of the halogen site) in apatite, amphibole, and biotite from martian meteorites. For apatite analyses where OH was not directly measured, it was calculated assuming 1–F–Cl = OH. For the amphiboles that do not commonly have oxy-components in the O(3) site (NWA 5790 and MIL nakhlites), OH was not directly measured, so we assumed that 2–F–Cl = OH, whereas kaersutite and Ti-biotite within Chassigny were only plotted if OH was measured directly because of the likelihood of O2− substitution in the O(3) site. To plot within the F-Cl-OH ternary, we normalized the molar sums of F, Cl, and OH from the SIMS analyses to one. a Apatites from shergottites and regolith breccia NWA 7034, b Apatites from chassignites, c Amphiboles and biotite from chassignites and nakhlites, and d Apatites from nakhlites. Data within these plots are compiled from McCubbin et al. (2013, 2016)

16.2.2 Nakhlites

The nakhlite meteorites are clinopyroxenite cumulates that consist of ~19 separate stones, believed to have comprised a single cumulate pile or thick lava flow on Mars. As discussed above, the bulk rock halogen abundances of cumulate rocks cannot be used to constrain the halogen abundances of their parental magmas because their bulk compositions do not represent magmatic liquids. However, the halogen mineralogy in the whole rock can be used to understand the role that halogens have played during both high-T and low-T processes, depending on the textural context in which the minerals occur. Halogen-bearing mineral phases have been shown to occur in three distinct petrographic contexts in nakhlites: (1) discrete phases within partially crystallized, olivine- or clinopyroxene-hosted melt inclusions; (2) igneous textured apatite that are interstitial to the cumulus olivine and clinopyroxene; and (3) secondary, low-temperature, hydrothermal- or aqueous-alteration products purported to be of martian origin. The partially crystallized melt inclusions host a number of halogen-bearing phases, including apatite, Cl-rich amphibole, and Cl-bearing jarosite (Treiman 1986, 1990, 1993; Harvey and McSween 1992; Aoudjehane et al. 2006; Sautter et al. 2006; McCubbin et al. 2009, 2013). However, not all of these phases are present in every inclusion, and some of the melt-inclusion minerals are specific to individual meteorites. Compositional information relating to the amphibole and apatite phases found in nakhlite melt inclusions is shown in Fig. 16.3. Halogen analyses of mineral separates following crushing also revealed elevations in halogen abundances, and are thought to represent components released from fluid inclusions (Cartwright et al. 2013). Apatite is the only primary, igneous, halogen-bearing mineral phase that occurs interstitial to the cumulus clinopyroxene and olivine. The intercumulus apatite compositions are displayed in Fig. 16.3. The nakhlites also have low-temperature, halogen-bearing, alteration phases that are likely martian in origin. These minerals include Cl-poor smectite, Cl-bearing iddingsite, and halite (Bridges and Grady 1999, 2000; Bridges et al. 2001; Treiman 2005; Bridges and Schwenzer 2012).

16.2.3 Chassignites

The three chassignite meteorites (Chassigny, NWA 2737, and NWA 8694) are dunites, and may originate from the same cumulate pile as the nakhlites. Halogen-bearing mineral phases occur in two distinct petrographic contexts in chassignites: (1) as discrete phases within crystallized, olivine-hosted melt inclusions; and (2) igneous textured apatite that is interstitial to the cumulus olivine. The crystallized melt inclusions have mineral assemblages that include kaersutitic amphibole, Ti-rich biotite, apatite, and Cl-OH-bearing shock-vitrified plagioclase (maskelynite) (Johnson et al. 1991; Watson et al. 1994; McCubbin and Nekvasil 2008; McCubbin et al. 2010, 2013). The amphibole, biotite, and apatite compositions from the olivine-hosted melt inclusions in the chassignites are shown in Fig. 16.3. The intercumulus regions of the chassignite meteorites host only apatite as a volatile-bearing mineral phase. The compositions of intercumulus apatites, shown in Fig. 16.3, have higher abundances of Cl than the apatites hosted by melt inclusions. This observation is reported to be the result of interactions between intercumulus melt in the chassignites and Cl-rich crustal fluids on Mars at the time of their formation (McCubbin and Nekvasil 2008; McCubbin et al. 2013).

16.2.4 Orthopyroxenite (Allan Hills 84001)

The single orthopyroxenite meteorite—Allan Hills 84001 (ALH 84001), has been shown to contain apatite phases that are interstitial to the cumulus orthopyroxenes (Boctor et al. 2003; Greenwood et al. 2008). A single study measured the light halogen abundances of these apatites, showing that they are Cl-rich with 3.42 wt.% Cl and 1.29 wt.% F (Boctor et al. 2003). Hydroxyl (OH) abundances have also been determined by secondary ion mass spectrometry (SIMS), yielding approximately 0.08–0.22 wt.% H2O (Boctor et al. 2003; Greenwood et al. 2008). As OH is a standard substitution site for Cl and other halogens, these data indicate that the apatites are predominantly halogen-rich, and similar in composition to apatites in the nakhlites and chassignites.

16.2.5 Regolith Breccia

Until recently, we had no non-igneous martian meteorites within our inventory, and thus no samples that could represent the crust. However, in 2011, a martian regolith breccia—NWA 7034—was recovered, and since this discovery, at least eight additional paired stones have been reported. This sample-set is considered to be the most representative of martian crust to-date. A bulk Cl content of 2200 ppm was determined for NWA 7034, which is similar to values reported for the martian crust from gamma ray spectroscopy, ranging from 350–4200 ppm (Taylor et al. 2010; Taylor 2013). The only halogen-rich phase within this regolith breccia is Cl-rich apatite, which is hosted in three distinct petrographic contexts: (1) as mineral clasts within the bulk matrix of the breccia (Agee et al. 2013; Humayun et al. 2013); (2) within igneous clasts that appear to be primary crystallization products (Santos et al. 2015); and (3) as submicron apatites that constitute the thermally annealed granoblastic groundmass of the breccia (Muttik et al. 2014). Apatite compositions from each of the petrographic domains are available in Fig. 16.3. The occurrence of these phases is important because apatite has been shown to account for the entire Cl budget of NWA 7034, so it provides us with at least one data point on the mineralogical hosts for Cl in the martian regolith. Although bulk rock F abundances have not been determined for NWA 7034, the abundance is estimated to be ~316 ppm from the average F/Cl ratio of the apatites throughout NWA 7034 (McCubbin et al. 2016), and the bulk rock abundance of Cl, which can be entirely attributed to Cl in apatite.

16.3 Significance of Halogens in Martian Meteorites

The study of halogens within martian meteorites has important implications for igneous, atmospheric, and aqueous processes on Mars, as well as representing an important analogue with which to quantify the budget of halogens in the martian interior. Some studies have highlighted the importance of Cl in lowering the basalt liquidus, which acts to reduce pressure or temperature parameters for crystallization and melting in the martian interior (Filiberto and Treiman 2009a, b; see also Webster et al. 2018). They suggested that Cl may have had a larger impact than H2O on basalt formation within the martian mantle, and thus the formation of martian rocks and soils, which may explain the elevated halogen concentrations observed in such materials compared to terrestrial basalts (see below). Later studies found that many martian meteorites have been affected by contamination from martian crustal material and, therefore, could not be used to reliably estimate volatile abundances in the martian mantle (McCubbin et al. 2016; Williams et al. 2016). Recently constrained volatile abundances show that the source region for martian meteorites is similar to the terrestrial Mid-Ocean-Ridge Mantle source, and melting of this source would produce magmas with roughly equivalent abundances of H2O and Cl (Filiberto et al. 2016; McCubbin et al. 2016).

The Cl and F abundances of the martian mantle and crust have been estimated and modeled from shergottites and regolith breccia samples, respectively, in a number of studies through methods that combine halogen mineral chemistry, bulk rock abundances of F, bulk rock abundances of incompatible-refractory-lithophile elements, and mineral-melt partitioning relationships (McCubbin et al. 2012, 2016; Gross et al. 2013; Filiberto et al. 2016). These models first determine the pre-degassing abundances of Cl in the parental liquids of the shergottites, and then ratio the F and Cl values for the parental liquids to the bulk rock abundances of incompatible-refractory-lithophile elements. These are used to calculate the F and Cl abundances of the source, where the abundances of incompatible refractory lithophile elements for bulk silicate Mars are constrained relatively well based on their chondritic abundances (Dreibus and Wanke 1985; McDonough and Sun 1995; Taylor 2013). The abundances of refractory, incompatible, lithophile elements are useful for estimating the abundances of volatile, incompatible, lithophile elements because they both exhibit similar behavior during magmatic processes. Based on these methods, it is estimated that the depleted shergottite source has 1.6–4.2 ppm Cl and 1.0–1.6 ppm F, and that the enriched shergottite source has 12–23 ppm Cl and 3.6–5.4 ppm F (McCubbin et al. 2016). In addition, the martian crust has 450 ppm Cl and 106 ppm F (McCubbin et al. 2016). Combined, these estimates imply that the bulk silicate composition of Mars is approximately 44 ppm Cl and 10 ppm F (McCubbin et al. 2016).

Recent work by Williams et al. (2016) and Sharp et al. (2016) assessed the Cl isotopic abundance in a number of martian meteorites and determined a range in 37Cl, where the martian mantle (olivine-phyric shergottites) is depleted at −3.8‰ (i.e., the difference in isotopic composition from standard mean ocean Cl in per mil, or parts per thousand), while the crust (NWA 7034) is enriched at up to 8.6‰. They conclude that the enrichment on the surface was likely caused by preferential loss of 35Cl to space, while the low 37Cl value likely relates to the primordial bulk composition of Mars, inherited during accretion.

A study by Cartwright et al. (2013) on nakhlites Nakhla, NWA 998, and MIL 03346 found evidence for a trapped fluid component enriched in Cl, which correlates with excess 40Ar (excess 40Ar = 40Armeasured40Arradiogenic, thus leaving the composition of the trapped fluid component), consistent with the martian atmosphere. The range of I/Cl and Br/Cl ratios for the mineral separates studied was similar to that observed on the Earth, which was surprising given the lack of crustal recycling processes, organic activity, and evident fluid activity on Mars. The authors also concluded that the halogen components are likely dispersed within minor phases in the nakhlites, and that the fluid observed had a low salinity, suggesting that it may have originated as a shallow sub-surface fluid within the nakhlite cumulate pile. A similar Cl-rich component was also observed in halogen analyses of shergottite RBT 04262 (Cartwright and Burgess 2011).

Earlier work by Bridges et al. (2000, 2001) studied the formation of alteration assemblages in the nakhlite suite and suggested that brines with a seawater-like composition could have percolated through the crust and deposited evaporite assemblages within the nakhlites under low temperature conditions. An additional source of fluid interaction could be through impact-induced heating. Later studies have suggested that impacts on the surface produced sufficient heat sources to drive hydrothermal activity at impact sites and the surrounding areas, leading to volatile formation in the crust, and alteration phases in the surface that include phyllosilicates and S-rich phases (Newsom 1980; Abramov and Kring 2005; Schwenzer and Kring 2009, 2010). A study by Changela and Bridges (2010) on alteration phases within nakhlites led the authors to conclude that the assemblages may result from the formation of an impact-induced ‘hydrothermal cell’, where a stratigraphic variation in alteration phase abundance was caused by fluid flow from the base to the top of the pile rather than top-to-bottom. Later work by Bridges and Schwenzer (2012) further modeled this hydrothermally driven fluid, determined that it was likely impact-generated, and that it had an alkaline pH. Additional studies have also modeled the flow of hydrothermal fluids following an impact on Mars and found clear processes of interaction with the martian surface (Abramov and Kring 2005; Schwenzer and Kring 2009).

An additional observation for martian meteorites is that elevated I/Cl ratios are observed in finds compared to falls, especially for those recovered from Antarctica (Cartwright et al. 2013). This is consistent with the findings of Langenauer and Krähenbühl (1993) who studied terrestrial halogen contamination in chondrites from Antarctica, and found some abundance patterns based on distance from the Antarctic coast. While falls are less likely to experience terrestrial weathering due to their shorter exposure time as a result of fast recovery in the field, some variation was observed in sub-samples of Nakhla (E, K, G—which relate to the museum/collection that they were provided from—Table 16.1). This was explained as either a result of contamination by salts from terrestrial brines prior to recovery, or the result of heterogeneous martian weathering (Dreibus et al. 2006). Here, Nakhla G is considered to be the most ‘pristine’ sample, and shows higher I/Cl ratios compared to E and K (Dreibus et al. 2006). While small-scale terrestrial contamination within Nakhla may be possible, it is also plausible that martian weathering is a factor. In fact, Greenwood (2008) suggested that Antarctic weathering may be a good analogue for martian weathering, where elevated halogen contents should be expected in martian weathering products.

16.4 In Situ Detections of Halogens on the Martian Surface

Halogens have been detected on the martian surface by all landed missions (Fig. 16.4, Table 16.2 ). Chlorine has been detected in all landed missions to date, whereas Br has only been measured in more recent missions, and F was recently measured by the Curiosity rover (Table 16.3). Iodine has not yet been measured on the martian surface, likely because of low abundances and instrument detection limits. Overall, surface rocks and soils have higher concentrations of halogens than martian meteorites. The type and the abundance of halogens and their speciation have important implications for ancient and modern environments on Mars. Here, we review the types of halogens that have been detected in situ by rovers and landers and the phases in which these halogens may reside. We discuss the variability and mobility of halogens on the surface and the processes that control this variability. Finally, we discuss the implications that the discovery of Cl in the form of perchlorate has for the modern habitability of the martian surface and the preservation of organics.
Table 16.2

Largest halogen abundances measured at each landing site

Mission

Landing site

Cl (wt.%)

Br (ppm)

Viking L1

Chryse Planitiae

0.9 + 1.5/−0.5*

(C-5 & C-13)

Viking L2

Utopia Planitiae

0.6 + 1.5/−0.5*

(U-2 & U-5)

Pathfinder

Ares Vallis

1.2 ± 0.3

(A-9)

MER—Opportunity

Meridiani Planum

2.13 ± 0.03

(Dorsal_new)

1232 ± 38 (Ellesmere_Barbeau)

MER—Spirit

Gusev Crater

2.62 ± 0.03§

(Uchben_Chiikbes_brush)

1543 ± 28§ (Temples_dwarf_asis)

Phoenix

Northern Plains

0.6 mM Cl; 2.6 mM ClO 4 − ‖

(Rosy Red)

<100 ppm

MSL—Curiosity

Gale Crater

3.44 ± 0.09

(Stephen_Raster2)

1978 ± 60 (Windjana_center_postDRT)

Targets of the measurements are in parentheses

*From XRFS data (Clark et al. 1982). Total uncertainty includes uncertainty with instrument precision, calibration, and soil matrix effects

From APXS (Foley et al. 2003). Error includes the statistical and laboratory combined error at 1σ

Values from the Planetary Data System. Absolute statistical 2σ errors are derived from peak area errors and do not include calibration uncertainties

§From APXS data (Gellert et al. 2006). Absolute statistical 2σ errors are derived from peak area errors and do not include calibration uncertainties

From WCL data (Hecht et al. 2009). Concentration error = ±20%

Br and I were not detected in solution, indicating they must be present in abundances below the WCL sensor’s limits of detection (Kounaves et al. 2010)

Fig. 16.4

Topographic map of Mars from the Mars Orbiter Laser Altimeter showing the landing sites of landers and rovers, important features, and locations discussed in the text.

Image credit MOLA Science Team

Table 16.3

Halogen abundances measured by APXS on Curiosity of drilled and scooped samples

Location

Sample

Lithology

Sol

Cl (wt.%)

Br (ppm)

Near landing site

Rocknest

Aeolian Sand

102

0.88 ± 0.05

65 ± 10

Yellowknife Bay

Cumberland

Mudstone

487

1.19 ± 0.04

65 ± 10

Kimberley

Windjana

Sandstone

704

0.57 ± 0.01

122 ± 5

Pahrump Hills

Confidence Hills

Mudstone

767

0.35 ± 0.01

38 ± 5

Pahrump Hills

Mojave2

Mudstone

888

0.46 ± 0.02

57 ± 5

Pahrump Hills

Telegraph Peak

Mudstone

922

0.34 ± 0.01

93 ± 5

Marias Pass

Buckskin

Mudstone

1065

0.28 ± 0.01

60 ± 5

Bridger Basin

Big Sky

Sandstone

1126

0.84 ± 0.02

300 ± 10

Bridger Basin

Greenhorn

Sandstone

1143

0.49 ± 0.01

219 ± 10

Bagnold Dunes

Gobabeb

Aeolian Sand

1223

0.50 ± 0.01

37 ± 5

Naukluft Plateau

Lubango

Sandstone

1326

0.32 ± 0.01

60 ± 5

Naukluft Plateau

Okoruso

Sandstone

1339

0.61 ± 0.01

115 ± 5

Murray Buttes

Oudam

Mudstone

1368

0.35 ± 0.01

24 ± 5

Measurements from rocks (i.e., mudstone and sandstone) were made on drilled fines prior to sieving, except for the measurement from Telegraph Peak, which is on drill tailings. Values are from the Planetary Data System

The NASA Viking Landers, VL1 and VL2, landed in 1976 in Chryse and Utopia Planitiae, respectively. The X ray Fluorescence Spectrometers (XRFS), mounted within the body of both landers, used a 55Fe and 109Cd sources and had the ability to detect fluorescent emissions of elements between the mass ranges of Mg and U, as long as they were above the limit of detection (Clark et al. 1977). The only halogen detected by the XRFS on VL1 and VL2 was Cl with a detection limit of ~0.3 wt.% (Clark and Baird 1973).

The NASA Mars Pathfinder mission featured the first rover, Sojourner, which landed in Ares Vallis in 1997. The alpha proton X-ray spectrometer (APXS) mounted on the back of the rover used a 244Cm source and operated in three separate modes (alpha backscattering, proton emission, and X-ray emission) to measure a different range of elements (from C to Zr) than the Viking XRFS (Rieder et al. 1997; Foley et al. 2003). A deployment mechanism allowed the instrument to investigate both rocks and soils (Rieder et al. 1997). Like the XRFS instruments on VL1 and VL2, the APXS on Pathfinder only detected Cl in martian rocks and soils at Ares Vallis.

The NASA Mars Exploration Rovers (MER), Opportunity and Spirit , landed in Meridiani Planum and Gusev crater, respectively, in 2004. Opportunity continues to rove the martian surface today. Both rovers had an updated APXS instrument mounted on the turret of a robotic arm. Upgrades from the Pathfinder-APXS included an improved X-ray detector, which gave higher energy resolution and sensitivity (Rieder et al. 2003; Gellert et al. 2006). The APXS instruments on both rovers detected Cl and made the first detections of Br on the martian surface (e.g., Rieder et al. 2004; Gellert et al. 2006).

The NASA Phoenix lander landed in 2008 in Vastitas Borealis in the northern plains of Mars (at 68.22°N). The Wet Chemistry Laboratory (WCL) on Phoenix had the capability to detect Cl, Br, and I anions in solution and indeed detected Cl in soil solutions (Hecht et al. 2009). The Thermal and Evolved Gas Analyzer (TEGA) on Phoenix showed that the Cl was present in oxychlorine compounds (Boynton et al. 2009; Hecht et al. 2009).

The NASA Mars Science Laboratory (MSL) Curiosity rover landed in Gale crater in 2012, and the mission is in its second Extended Mission at the time of this writing. An APXS instrument mounted on the turret of the arm has detected Cl and Br in nearly every measurement to date (e.g., Mangold et al. 2017; Thompson et al. 2016, 2017), and the Sample Analysis at Mars (SAM) instrument has confirmed the presence of oxychlorine compounds from evolved gas analysis (EGA) of most drilled rock samples and scooped sand samples measured to date (e.g., Glavin et al. 2013; Leshin et al. 2013; Archer et al. 2014, 2016; McAdam et al. 2014; Sutter et al. 2016, 2017; Stern et al. 2017). Curiosity has made the first in situ measurement of F by laser induced breakdown spectroscopy (LIBS) using the ChemCam instrument (Forni et al. 2015). This technique is currently being developed to quantify Cl abundances (Thomas et al. 2017).

16.4.1 Chlorine

Chlorine is the most abundant halogen detected on the martian surface by landers and rovers (Table 16.3 and Fig. 16.4). Early detections of Cl by the Viking landers and MER were generally proposed to be in the form of chloride salts (e.g., Clark and Baird 1979; Clark et al. 2005; Knoll et al. 2008). Enrichments in both Na and Cl in rinds on rocks in Meridiani measured by Opportunity suggest the presence of halite on rock surfaces, which may have formed by transient thin films of liquid water (Knoll et al. 2008). Curiosity is ascending the lower slopes of Mount Sharp in Gale crater to study progressively younger sedimentary strata (Fig. 16.5), which were deposited in lacustrine, fluvial, deltaic, and aeolian environments at ~3.6 Ga. Curiosity is currently investigating the Murray formation (Fig. 16.5b), which is dominated by lacustrine mudstone in the lower Murray formation in Pahrump Hills and Marias Pass, but is made up of lacustrine, fluvial, and aeolian deposits in the upper Murray formation south of the Murray Buttes. Association of Cl with Na and a greater abundance of Cl have been observed in the upper Murray formation compared to the lower Murray formation (Thomas et al. 2017). This indication of halite, along with the presence of Ca-sulfate minerals, hematite, and desiccation cracks, suggests that this mudstone was deposited in an evaporative lake environment (Stein et al. 2017; Vaniman et al. 2017). This is in contrast to the perennial lake environment inferred from finely laminated mudstone investigated at the Pahrump Hills and Marias Pass (Rampe et al. 2017).
Fig. 16.5

a View of Gale crater from a combination of data from three Mars orbiters. The crater is 154 km in diameter, and the landing ellipse is in black. Location of Fig. 16.5b is outlined in red. Image credit NASA/JPL-Caltech/ESA/DLR/FU Berlin/MSSS. b Curiosity rover traverse, showing the important waypoints through Sol 1500. At the time of writing, Curiosity is located at the gold star.

Image credit NASA/JPL-Caltech/Univ. of Arizona

Clark et al. (2005) noted that Cl on the martian surface could be present in oxidized compounds, including chlorites, chlorates, and perchlorates. Data from TEGA on Phoenix and SAM on Curiosity indeed indicate that much of the Cl detected in samples measured from the northern plains and Gale crater is in the form of oxychlorine compounds chlorate (ClO3 ) and/or perchlorate (ClO4 ) (e.g., Boynton et al. 2009; Archer et al. 2014; Sutter et al. 2016, 2017; Stern et al. 2017). The identification of oxychlorine compounds is based on O2 gas releases from ~100–600 °C in EGA data. The position of this release can help differentiate between chlorate and perchlorate and can help identify the cation associated with the oxychlorine (Sutter et al. 2015; Archer et al. 2016). The presence of perchlorate was inferred at the Viking landing sites based on the detection of chloromethane and dichloromethane in the Viking Gas Chromatograph Mass Spectrometer (Navarro-González et al. 2010). Oxychlorine compounds in modern aeolian sediments and ancient sedimentary rocks in Gale crater are generally much more abundant than in terrestrial soils (106–107 μg/kg vs. 10−1–105 μg/kg, respectively; Stern et al. 2017), though similar abundances have been identified in soils from the Atacama Desert (Catling et al. 2010). Wilson et al. (2016) and Stern et al. (2017) suggested that this discrepancy may give us insight into the mechanisms by which oxychlorine compounds form on Mars.

Oxychlorine compounds have been hypothesized to form by a variety of processes on Mars, including oxidation on mineral surfaces, radiolysis by ionizing radiation, and/or UV irradiation. Recent laboratory studies and in situ measurements on Mars point toward radiolysis by ionizing and/or UV radiation as an important process for forming and destroying oxychlorine compounds on Mars (e.g., Quinn et al. 2013; Carrier and Kounaves 2015; Wilson et al. 2016; Stern et al. 2017). Wilson et al. (2016) proposed a mechanism for the production of perchlorate through radiolysis of the martian surface by galactic cosmic rays from which Cl oxides sublimate into the atmosphere and form perchloric acid (HClO4). Perchloric acid is then deposited onto the surface and crystalizes to form perchlorate salts. The similarity in oxychlorine abundances in rocks from the early Hesperian period (~3.7–3.4 Ga) rocks and in modern sediments as measured by SAM in Gale crater suggests ionizing radiation may play a role in both the formation and destruction of oxychlorine compounds to maintain a steady surface concentration (Stern et al. 2017). UV radiation may also add to the concentration of oxychlorine compounds on the martian surface. This has been demonstrated by Carrier and Kounaves (2015), whom exposed grains of halite and silica to UV radiation resulting in the oxidation of Cl.

Mudstone sampled from the Yellowknife Bay sampling site in Gale crater and sandstone sampled from the Kimberley outcrop measured by SAM on Curiosity have light and highly variable δ37Cl values (Farley et al. 2016). The isotopic composition of Cl was measured from thermally evolved HCl. The δ37Cl values range from −1 ± 25‰ to −51 ± 5‰. The mineral phase in which the Cl resides has important implications for the reasons behind the isotopically light Cl. If the Cl is in oxychlorine compounds, then atmospheric chemical reactions may be responsible for the light Cl. Similar reactions create isotopically light Cl in Atacama Desert soils (e.g., Jackson et al. 2010). Alternatively, if the Cl is in chloride salts, then partial reduction of isotopically normal perchlorate could fractionate and produce isotopically light chloride. The CheMin X-ray diffractometer on Curiosity did not definitively detect oxychlorine compounds or chloride salts in any of these samples, suggesting they are below the instrument’s detection limits of ~1–2 wt.% (Blake et al. 2012). Although it is not certain whether the Cl in the evolved HCl was from an oxychlorine or chloride source, the low-temperature evolution of O2 with the concomitant release of HCl implicates oxychlorine (e.g., Glavin et al. 2013; Leshin et al. 2013; Archer et al. 2014; Ming et al. 2014).

Whatever its mechanism of formation, the discovery of oxychlorine in areas of the martian surface so-far- sampled demonstrates the paucity of liquid water since its deposition. Oxychlorine compounds are extremely soluble in liquid water and will persist in soils as long as there is no water to mobilize them (e.g., Catling et al. 2010). The samples measured by SAM are sourced from ~5–6 cm depth, so the discovery of oxychlorine compounds in these ancient rocks (~3.6 Ga) suggests either that the oxychlorine compounds are ancient or that late-stage diagenetic fluids have affected their concentrations and mobilized them from elsewhere in the crater. Indeed, in two samples drilled from mudstone at Yellowknife Bay, there was a marked difference between the oxychlorine abundance in the two samples, named Cumberland and John Klein, drilled only a few meters apart. The ClO4 abundances in the Cumberland drill sample were an order of magnitude greater than those in the John Klein sample (10.5 ± 4.5 vs. 0.87 ± 0.39 g/kg, respectively; e.g., Ming et al. 2014; Stern et al. 2017). Images of the John Klein drill hole (Fig. 16.6) and geochemical measurements of the walls by ChemCam showed hairline fractures filled with Ca-sulfate cement (Vaniman et al. 2014), suggesting that later stage fluids removed a majority of the oxychlorine compounds from that part of the mudstone.
Fig. 16.6

Mars Hand Lens Imager (MAHLI) images of the John Klein and Cumberland drill holes, taken on sols 270 and 293, respectively. For scale, the drill holes are 1.6 cm in diameter. Note the fine, bright veins in the walls of the John Klein drill hole (highlighted in yellow) and the relative paucity of veins in the walls of the Cumberland drill hole. The late-stage fluids in these fractures could have mobilized oxychlorine compounds in the mudstone on a local scale.

Image credit NASA/JPL-Caltech/MSSS

16.4.2 Bromine

Bromine is present in almost all APXS measurements by MER (Opportunity and Spirit) and MSL (Curiosity ). Bromine concentrations on Mars are much higher than on Earth, on the order of 1–10% of Cl concentrations, whereas Br concentrations in seawater on Earth are 0.35% of Cl (e.g., Gellert et al. 2006; Marion et al. 2009). On Earth, Br typically substitutes for Cl in halite, rather than precipitating as a discrete bromide mineral. To evaluate the phases in which Br may precipitate on Mars, Marion et al. (2009) investigated the partitioning of Br in chloride minerals in the Burns formation (a 7 m exposed stratigraphic section in Meridiani Planum) using an equilibrium chemical thermodynamic model, FREZCHEM. They found that the mineral bischofite (MgCl2 · 6H2O) may be a more important sink for Br than halite on Mars. Jarosite ((K,Na,H3O)Fe3(SO4)2(OH)6) has been unequivocally identified by rovers in Meridiani Planum and Gale crater (e.g., Klingelhöfer et al. 2004; Rampe et al. 2017). Laboratory experiments of jarosite precipitation in Br- and Cl-bearing solutions demonstrate that jarosite may be another mineral host for Br (Zhao et al. 2014). Others have speculated that Br is present in oxidized forms, as bromates and perbromates (Clark et al. 2005).

Many studies have noted extremely variable Br concentrations, particularly in Meridiani and Gusev. Concentrations of Br at the surface of Meridiani and Gusev soils are typically less than 50 ppm, but these concentrations are elevated by factors of 2–30 in subsurface soils, grains armoring sedimentary bedforms, and low-lying rocks (Yen et al. 2005). Enrichments in Br and Cl are decoupled from one another in these soil profiles, where APXS measurements show a relative loss of Br compared to Cl at the surface (Clark et al. 2005; Karunatillake et al. 2013). Initially, this variability was attributed to the presence of liquid water and the high solubility of bromide in solution. It was hypothesized that thin films of liquid water could form from the sublimation of frost, and cycles of diurnal or seasonal thin films of liquid water could mobilize Br and concentrate it in the subsurface and local depressions (Yen et al. 2005 ). Bromine concentrations, however, show no obvious relationship with other highly soluble salts (Clark et al. 2005). More recent studies have suggested that UV photolysis and the subsequent oxidation of Br may explain the lower concentrations of Br in soil surfaces (Karunatillake et al. 2013). In this scenario, Br is converted to gas phases (e.g., BrO) by UV photolysis. In fact, depletion of Br from salt flats in Abu Dhabi and from terrestrial polar sea ice is observed and attributed to UV photochemical reactions and the formation of Br in the vapor phase (Simpson et al. 2007; Wood and Sanford 2007).

16.4.3 Fluorine

Fluorine was recently measured in Gale crater with ChemCam (Curiosity ), marking the first in situ measurement of F on Mars. Atomic or ionic emission lines from F in laser induced breakdown spectra (LIBS) from MSL-ChemCam have a detection limit of ~5 wt.% because of atmospheric absorption, laser coupling, and detector sensitivity (Forni et al. 2015). However, F can be detected by ChemCam with a detection limit of ~0.2 wt.% by observing CaF molecular bands. As of early 2017, F had been identified in 600+ targets (Forni et al. 2017). The highest concentration was found in the Alvord Mountain target at 10 wt.% F in the lower Murray formation in a location with a complex series of veins called Garden City (Figs. 16.5b, 16.7). Fluorine has been identified in a variety of targets, suggesting multiple sources and formation mechanisms. It has been identified (1) in conglomerates where the F-bearing phase is an aluminosilicate, (2) in conglomerates where the F-bearing phase is high in Ca (e.g., fluorapatite or fluorite), (3) associated with Si at the base of the Pahrump Hills outcrop, (4) in dark veins in Garden city and fracture fills in aeolian sandstone where the F-bearing phase is high in Ca, and (5) in correlation with Ca near the unconformity between lacustrine mudstone and aeolian sandstone (Forni et al. 2015, 2016, 2017). The detection of F, associated with Al and/or Si, may indicate that F is in the phyllosilicates. The association of F with Ca suggests the presence of fluorapatite and/or fluorite. Fluorapatite was detected by the CheMin X-ray diffractometer in abundances of ~1–2 wt.% in mudstone samples from the Pahrump Hills (Rampe et al. 2017). The prevalence of Ca-F minerals in veins and along unconformities suggests they precipitated from late-stage, low-temperature diagenetic fluids (Forni et al. 2015, 2016, 2017).
Fig. 16.7

Curiosity Mastcam image of complex series of veins at the Garden City site.

Image credit NASA/JPL-Caltech/MSSS

16.5 Orbital Detections of Halogens on the Martian Surface

Halogens have been detected on a global scale by orbiting spectrometers, providing clues to past and present geologic environments and processes.

16.5.1 Chlorine

The first measurements of Cl from orbit were made from the 2001 NASA Mars Odyssey Gamma Ray Spectrometer (GRS). Keller at al. (2006) determined a mean concentration of 0.49 wt.% Cl from the summation of the GRS spectra collected over the planet. Consistent with lander measurements and the meteorite isotope work performed at the time (Rao et al. 2002), Cl is enriched significantly within the upper few tens of centimeters of the surface relative to the martian meteorites and the estimate for the bulk composition of the planet. However, Cl is not homogenously distributed over the martian surface and varies by a factor of ~4. Only regional variations can be investigated by GRS due to the instrument’s relatively large footprint (~600 km diameter). Keller et al. (2006) observed that the Medusae Fossae formation west of the Tharsis volcanic complex shows significantly elevated Cl, whereas distinctly low Cl values are observed in the southern highlands and in a region north of Syrtis Major extending into Utopia Planitia. Furthermore, measurable differences from the global mean are observed around the outflow channels of Chryse and Acidalia Planitiae and Arabia Terra. Chlorine appears to be positively linked with H and negatively associated with Si and thermal inertia (Keller et al. 2006).

Genetic relationships between Cl enrichment and depletion are likely due to many factors. The strong spatial overlap between the Cl-rich region west of Tharsis and the previously mapped Medusae Fossae formation suggest enrichments through volcanic outgassing given the proximity to Tharsis. The denudation of volcanic ignimbrite deposits enriched in Cl through reactions of acid-fog or acid precipitation, in addition to possible aqueous activity, could result in Cl enrichment in the region (Keller et al. 2006). However, movement of Cl in ground or surface water could have depleted Cl in some regions, such as the ancient southern highlands, through leaching and erosion and deposition in other regions through evaporitic processes (e.g., Kargel 2004). Because of weak correlations with surface dust, Cl distribution is also likely the result of aeolian processes resulting in the deposition of Cl-rich fines. The Cl-bearing species in the dust have not been determined because the physical nature of the fine-grained surface dust makes it difficult to decipher the composition from orbiting visible/near-infrared (VNIR) and mid-infrared (MIR) spectrometers.

16.5.2 Halite-Bearing Materials

The Thermal Emission Imaging System (THEMIS) onboard NASA’s Mars Odyssey provided the first orbital evidence for chloride salts on the surface (Osterloo et al. 2008). The chloride salt deposits were discovered based on their spectral distinctiveness in false-color THEMIS images. The materials exhibit featureless sloping emissivity spectra in THEMIS (~672–1475 cm−1 or ~14.88–6.78 μm) and NASA’s Mars Global Surveyor (MGS) Thermal Emission Spectrometer (TES) data (~300–1300 cm−1 or ~33.33–7.70 μm). The observed spectral slope is due to an erroneous assumption of unit emissivity (i.e., maximum emissivity of one over the wavelength range) in the conversion of measured radiance to emissivity (e.g., Ruff et al. 1997; Osterloo et al. 2008, 2010). Very few geological materials exhibit relatively featureless spectra over this region as well as non-unit emissivity with the exception of some anhydrous chloride salts, such as halite (NaCl). VNIR spectral data obtained by the NASA Mars Reconnaissance Orbiter (MRO) Compact Reconnaissance Imaging Spectrometer for Mars (CRISM) (Murchie et al. 2009) as well as the European Space Agency (ESA) Mars Express Observatoire pour la Minéralogie, lEau, les Glaces et l’Activité (OMEGA) (Ruesch et al. 2012) indicate that the chloride salts are distinctively brighter when compared to the typical background soils. Ratio reflectance data over the deposits exhibit featureless red slopes over the 1–2.5 μm region and an inverted 3 μm feature, indicating that these deposits are desiccated compared to the surrounding regolith. Although halite does not contain strong spectral features over the VNIR-MIR wavelength range, a combination of laboratory (Jensen and Glotch 2011), theoretical studies (Glotch et al. 2016), and geological arguments (e.g., Osterloo et al. 2008, 2010; Glotch et al. 2010), strongly suggest that it is the most likely candidate for the chloride salt phase present on Mars. Halite abundance is constrained to 5–20 wt.% for most deposits (Jensen and Glotch 2011). Furthermore, although phyllosilicates have been observed in close proximity to some deposits (Glotch et al. 2010; Ruesch et al. 2012), there appear to be no hydrated or additional evaporite phases intermixed with the chloride salts.

Based on THEMIS detections, the number of halite-bearing deposits ranges into the hundreds, dispersed primarily in the most ancient (i.e., Noachian-aged) terrains in the southern highlands (Fig. 16.8). The deposits occur in a variety of geologic settings but commonly occur in local depressions on plains units, including filling craters and channels (Osterloo et al. 2010; Osterloo and Hynek 2015; Hynek et al. 2015). Based on these observations, the most likely depositional scenario for halogens includes precipitation from evaporating surface waters in closed basins. Cross cutting relationships with phyllosilicates in close proximity to the deposits indicate that the halite-bearing materials are likely the products of late stage aqueous activity and are unlikely to be geologically related to the older phyllosilicates (e.g., Glotch et al. 2010). The lack of observed sulfates, carbonates, silica, phyllosilicates, and other expected authigenic phases, intermixed with the chloride salts, are problematic for drawing a clear link to a long-lived lacustrine depositional setting.
Fig. 16.8

Topographic map of Mars from the Mars Orbiter Laser Altimeter showing the locations where halide minerals have been detected by thermal infrared spectroscopy (black squares), oxychlorine phases have been detected by both near-infrared spectroscopy on orbiters and evolved gas analysis on Phoenix and Curiosity (white squares), and oxychlorine phases have been inferred at the VL1 and VL2 landing sites (yellow squares).

Image credit MOLA Science Team

The geologic setting, geomorphological observations, and relative ages of the halite-bearing surfaces indicate that they may represent the last surge of major surface water activity on Mars (Osterloo and Hynek 2015). Halite crystals have been shown to provide information about the precursory liquid compositions entrapped in inclusions and clues to the depositional environment and climate at the time of formation (Lowenstein et al. 1999; Roedder 1984). As such, these materials likely hold important clues regarding the paleoclimate and environment of this intriguing and dynamic time period of martian geologic history. Furthermore, terrestrial salt deposits contain microbial fossils (e.g., Huval and Vreeland 1991), cellulose fibers (e.g., Griffith et al. 2008), and microscopic prokaryotes, including both haloarchaea and halobacteria biomarkers (e.g., Fredrickson et al. 1997; McGenity et al. 2000; Barbieri et al. 2006; Schubert et al. 2009). In addition, fluid inclusions in halite have been shown to be a favorable refuge for the short to long term (possibly hundreds of millions of years) survival of halophilic microorganisms (e.g., Norton and Grant 1988; Javor 1989; Norton et al. 1993; Vreeland et al. 2000; Fish et al. 2002; Satterfield et al. 2005; Schubert et al. 2009). One of the primary goals of martian exploration is to determine whether life ever existed on the planet. Considering the ability of salts to both preserve and harbor microbial life, these halite deposits should be considered for future sample return missions.

16.5.3 Perchlorates

A limited number of detections of perchlorate and chlorate salts have been made from orbit (Fig. 16.8). Like halite, anhydrous oxychlorine salts do not have any strong features in the VNIR (Hanley et al. 2015). VNIR spectra of hydrated oxychlorine salts have many bands near 1.4 and 1.9 μm to allow their detection by orbital spectroscopy (Hanley et al. 2015). However, the VNIR spectra of hydrated oxychlorine salts are similar to some hydrated sulfate salts (Hanley et al. 2015). VNIR observations from OMEGA and CRISM demonstrate that hydrated sulfates are widespread across the surface of Mars (e.g., Bibring et al. 2006; Carter et al. 2013). As such, some detections of hydrated sulfates may actually be hydrated oxychlorine salts and vice versa (Hanley et al. 2015).

Data from MRO-CRISM provide evidence for perchlorates and chlorates over recurring slope lineae (RSL) in some locales (Ojha et al. 2015). RSL are narrow, low-reflectance features that form on present-day Mars (McEwen et al. 2011). These features are thought to form from either transient flow of liquid water or from granular flow of material (McEwen et al. 2011, 2014; Chevrier and Rivera-Valentin 2012). RSL extend incrementally downslope on steep, relatively warm slopes, fade when inactive, and reappear annually over multiple Mars years (Fig. 16.9) (McEwen et al. 2011, 2014; Ojha et al. 2014). Generally, RSL range from a few meters in length (<5 m), down to the detection limit for the High Resolution Imaging Science Experiment (HiRISE) camera (~0.25 m/pixel). The coarse spatial sampling of CRISM (~18 m/pixel), compared to HiRISE, limit the number of RSL that can be investigated with VNIR spectroscopy. However, Ojha et al. (2015) identified a few locations that were sufficient in size and/or number to undertake compositional studies using CRISM.
Fig. 16.9

Image of RSL in Palikir crater taken by HiRISE. RSL appear as dark streaks on northwest-facing slopes (highlighted by white arrows).

Image credit NASA/JPL/University of Arizona

CRISM pixels measured over a densely packed area of RSL in Palikir crater (Fig. 16.9) contained absorption features consistent with the presence of oxychlorine compounds. Pixels closest to the RSL exhibited features near ~1.48, 1.91, and ~3 μm, whereas pixels farther away only had features at ~1.91 and ~3 μm. The ~1.4 μm feature weakens with dehydration and disappears more rapidly than the 1.91 and 3 μm absorption bands suggesting higher hydration states in areas closest to the center of the RSL. Ojha et al. (2015) found that the wavelength position of the observed 1.4 μm absorption band is longer than is typical for perchlorates, suggesting the presence of an additional mineral. The absorptions observed are too narrow to be explained by liquid water and instead are hypothesized to be consistent with hydrated salts. A linear spectral mixture of martian soil with magnesium perchlorate, chlorate, and chloride salts provides the closest spectral match to what was observed in Palikir crater with CRISM (Ojha et al. 2015). A similar mixture of Mg-oxychlorine and chloride minerals was inferred from CRISM spectra of RSL in Hale crater (Ojha et al. 2015). Further evidence of hydrated salts comes from RSL on two central peaks in Horowitz crater, which show spectra best modeled by a linear mixture of martian soil and sodium perchlorate. RSL are abundant in Coprates Chasma, but spectra over these features only exhibited 1.9 μm absorptions, which limited assignment of a particular salt mineralogy (Ojha et al. 2015). Another notable detection of perchlorate from orbit using CRISM, although not from a region with RSL, is in dune sediments and sublimation tills surrounding the martian North Polar Cap (Massé et al. 2012). This perchlorate is associated with gypsum and may form by sublimation of the polar ice cap.

The origin of the water forming the RSL is not well understood (McEwen et al. 2011, 2014; Ojha et al. 2014). However, despite the uncertainty in formation mechanisms, hydrated salts likely play a role. Sodium perchlorate can lower the freezing point of water by up to 40 K, whereas magnesium perchlorate and magnesium chlorate can depress the freeing point even more by up to 70 K (Chevrier et al. 2009; Hanley et al. 2012). The presence of perchlorates and chlorates support the hypothesis that the RSL are formed from brine-enriched liquids on contemporary Mars. In the Atacama Desert on Earth, deliquescence of hygroscopic salts is a refuge for microbial communities (Davila et al. 2008, 2013b) and halophylic prokaryotes (Aharonson et al. 2014). This would suggest that if the RSL provide transiently wet conditions near the surface, they may have astrobiological significance, although the water activity in perchlorate solutions may be too low to support known terrestrial life (Rummel et al. 2014).

16.6 Oxychlorine Compounds and Implications for the Habitability and Exploration of the Martian Surface

The presence of oxychlorine compounds on Mars has important implications for the detection and preservation of organic molecules, habitability of the modern martian surface, and for eventual human exploration of the martian surface. Although perchlorate is a strong oxidizing agent, it is typically inert under Mars ambient conditions (Catling et al. 2010; Kounaves et al. 2014). The pathways by which perchlorate is generated, however, produce intermediates (e.g., hypochlorite (ClO), chlorite (ClO2 ), chlorate (ClO3 ), and their associated radicals) that could oxidize all but the most refractory organic molecules (Kounaves et al. 2014). Trace abundances of native martian organic molecules have been detected on the martian surface by SAM (e.g., Freissinet et al. 2015), but it is important to consider the role oxychlorine compounds may have played in forming the assemblage of organic molecules on the surface today. By comparison, the presence of oxychlorine compounds may enhance the microbial habitability of the modern surface by providing liquid water. Perchlorate can depress the freezing point of water to between 206 and 273 K, depending on salt concentration and relative humidity (Chevrier et al. 2009; Gough et al. 2011; Hanley et al. 2012).

Perchlorate has been discussed as both a resource and a potential hazard to the human exploration of Mars (e.g., Davila et al. 2013a). It could provide oxygen for life-support and fuel, yet it could be a hazard to humans if ingested, because it can inhibit the uptake of I ions by the thyroid, therefore resulting in thyroid hyperplasia, goiter, and reduced metabolism (e.g., Smith 2006; Davila et al. 2013a). Astronauts could be exposed to perchlorate through inhalation of dust, consumption of food grown in soils containing perchlorate, and in drinking water. The reference dose for perchlorate equates to 24.5 μg/l in drinking water (Brown and Gu 2006; Davila et al. 2013a). Davila et al. (2013a) suggest a biochemical pathway to remove O2 from perchlorate in the soil, where enzymes separate ClO4 to O2(g) and Cl. They suggest that this process would not violate planetary protection requirements because terrestrial microbes would not be necessary, and the resulting perchlorate-free soil could be used to grow crops.

16.7 Potential Measurements by Future Mars Missions

Many future orbital and landed missions to Mars are planned over the next 3–5 years, and this phase of Mars exploration will involve countries from around the world. NASA has two lander missions planned: InSight, which will have geophysical instruments to examine the interior, and the Mars 2020 rover, which will continue in Curiosity’s tire treads to examine the habitability of the martian surface. Both ESA and the Russian Federal Space Agency will collaborate to send the ExoMars rover in 2020 to search for biomarkers. Three additional missions are also planned for 2020: India plans to send a lander and a rover to follow on to its Mars Orbiter Mission; China plans to send an orbiter, a lander, and a small rover; and the United Arab Emirates plans to send an orbiter.

A few of the instruments on ExoMars and Mars 2020 will have the capability to identify some halogens or halogen-bearing minerals in situ. The MicrOmega infrared spectrometer will cover the spectral range from 0.5–3.5 μm, and the Infrared Spectrometer for ExoMars (ISEM) will cover the spectral range from 1.15–3.3 μm, both allowing for the detection of oxychlorine salts. Both ExoMars and Mars 2020 will have Raman spectrometers (the Raman Laser Spectrometer (RLS) on ExoMars, and the Scanning Habitable Environments with Raman and Luminescence for Organics and Chemicals (SHERLOC) and SuperCam instruments on Mars 2020). Although halite does not have any first-order Raman modes (e.g., Kieffer 1979), oxychlorine compounds are detectable by Raman (e.g., Wu et al. 2015). Like ChemCam on Curiosity , SuperCam on Mars 2020 will have a LIBS instrument capable of detecting Cl and F. There will not be an APXS instrument on Mars 2020, but the Planetary Instrument for X-ray Lithochemistry (PIXL) will determine bulk chemistry though X-ray fluorescence. PIXL will be able to detect Cl and Br directly, and scanning the beam may be able to put halogen geochemistry into a spatial context (e.g., identify whether halogens are present in veins, individual grains, or cements).

In addition to its science payload, Mars 2020 will have the ability to collect samples and cache them on the surface for a future sample return mission. With returned samples, the halogen geochemistry and speciation of these elements can be determined with greater precision using the high-powered instruments we have in terrestrial labs. As a word of caution: many of the mineral phases in which halogens reside are hydrated and can be hygroscopic. If the temperature and relative humidity in the sample tubes do not mimic those of the near subsurface, the halogen-bearing minerals could experience phase changes so that these minerals, when they’re analyzed on Earth, are not the same as those actually present on Mars. Curating the returned samples under cold temperatures will also be important for preventing further phase changes.

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Copyright information

© Springer International Publishing AG 2018

Authors and Affiliations

  • Elizabeth B. Rampe
    • 1
  • Julia A. Cartwright
    • 2
  • Francis M. McCubbin
    • 1
  • Mikki M. Osterloo
    • 3
  1. 1.NASA Johnson Space CenterHoustonUSA
  2. 2.Department of Geological Sciences, Bevill BuildingThe University of AlabamaTuscaloosaUSA
  3. 3.Laboratory for Atmospheric and Space PhysicsUniversity of ColoradoBoulderUSA

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