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Chemie der Erde 73 (2013) 401– 420

Contents lists available at ScienceDirect

Chemie der Erde

jou rn al homepage: www.elsev ier .de /chemer

nvited review

he bulk composition of Mars

. Jeffrey Taylor ∗

awaii Institute of Geophysics and Planetology, School of Ocean and Earth Science and Technology, University of Hawaii, Honolulu, HI 96822, United States

r t i c l e i n f o

rticle history:eceived 16 July 2013ccepted 11 September 2013

eywords:ars

ompositionlanet formationosmochemistryeochemistryerrestrial planets

a b s t r a c t

An accurate assessment of the bulk chemical composition of Mars is fundamental to understanding plan-etary accretion, differentiation, mantle evolution, the nature of the igneous parent rocks that were alteredto produce sediments on Mars, and the initial concentrations of volatiles such as H, Cl and S, importantconstituents of the Martian surface. This paper reviews the three main approaches that have been usedto estimate the bulk chemical composition of Mars: geochemical/cosmochemical, isotopic, and geophys-ical. The standard model is one developed by Wänke and Dreibus in a series of papers, which is based oncompositions of Martian meteorites. Since their groundbreaking work, substantial amounts of data havebecome available to allow a reassessment of the composition of Mars from elemental data, includingtests of the basic assumptions in the geochemical models. The results adjust some of the concentrationsin the Wänke–Dreibus model, but in general confirm its accuracy. Bulk silicate Mars has roughly uniformdepletion of moderately volatile elements such as K (0.6 × CI), and strong depletion of highly volatile ele-ments (e.g., Tl). The highly volatile elements are within uncertainties uniformly depleted at about 0.06 CIabundances. The highly volatile chalcophile elements are likewise roughly uniformly depleted, but withmore scatter, with normalized abundances of 0.03 CI. Bulk planetary H2O is much higher than estimatedpreviously: it appears to be slightly less than in Earth, but D/H is similar in Earth and Mars, indicating acommon source of water-bearing material in the inner solar system. K/Th ranges from ∼3000 to ∼5000

among the terrestrial planets, a small range compared to CI chondrites (19,000). FeO varies throughoutthe inner solar system: ∼3 wt% in Mercury, 8 wt% in Earth and Venus, and 18 wt% in Mars. These dif-ferences can be produced by varying oxidation conditions, hence do not suggest the terrestrial planetswere formed from fundamentally different materials. The broad chemical similarities among the terres-trial planets indicate substantial mixing throughout the inner solar system during planet formation, assuggested by dynamical models.

© 2013 Elsevier GmbH. All rights reserved.

ontents

1. Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4022. Approaches to estimating bulk composition . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 402

2.1. Models based on geochemistry and nebular components . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4022.1.1. Wänke and Dreibus model . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4022.1.2. Morgan and Anders model . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 404

2.2. Estimates based on isotopic composition of Martian meteorites . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4042.3. Estimates based on geophysical properties of Mars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4052.4. Summary of the models . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 405

3. Datasets . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4054. Complications . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 406

4.1. Is the Martian surface composition representative of the entire crust?. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4064.1.1. Highlands megaregolith . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 406

4.1.2. Lava flows . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

4.1.3. Element ratios . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .4.2. Heterogeneity of the Martian mantle . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

∗ Tel.: +1 808 956 3899; fax: +1 808 956 6322.E-mail address: gjtaylor@higp.hawaii.edu

009-2819/$ – see front matter © 2013 Elsevier GmbH. All rights reserved.ttp://dx.doi.org/10.1016/j.chemer.2013.09.006

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 406 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 406. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 406

402 G.J. Taylor / Chemie der Erde 73 (2013) 401– 420

5. A reassessment of Martian bulk composition . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4075.1. Estimating uncertainties . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4075.2. Refractory element abundances. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4085.3. FeO and MnO . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4095.4. Phosphorous . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4115.5. Moderately volatile elements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4115.6. Highly volatile elements, including halogens . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4125.7. Ni and Co . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4125.8. Strongly siderophile elements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4125.9. H2O and D/H . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 414

6. Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4156.1. Mars is rich in FeO . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4166.2. Depletion of volatile elements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4166.3. Water: abundance and source . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4166.4. Highly siderophile elements: implications for accretion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4176.5. Halogen concentration of the crust . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4176.6. The core . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4176.7. Comparing planet compositions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 417Acknowledgements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 418

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References . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

. Introduction

An accurate assessment of the bulk chemical compositionf Mars is fundamental to the entire geologic history of Mars,ncluding accretion, differentiation, aqueous alteration to produceediments, and the initial concentrations of important volatile ele-ents. For example, knowing the bulk composition in principle

llows us to understand crystallization and cumulate overturn inhe magma ocean (if there was one), partial melting to producehe crust through intrusion and extrusion of basaltic magmas, andhe formation of the distinctive source regions (e.g., enriched andepleted shergottites) of Martian meteorites. These processes, ofourse, make it tricky to extract the bulk composition from geo-hemical and geophysical data. An example of the complexity inust modeling magma ocean crystallization and cumulate overturnan be found in Elkins-Tanton et al. (2003, 2005). Fortunately, weave a solid database for the composition of the surface of Marsnd a good understanding of element behavior during petrologicrocessing. The database includes published meteorite analyses ofhe continuously expanding collection of Martian meteorites, andrbital and lander datasets.

This paper reviews models for Martian bulk composition andakes a complete reassessment in light of the substantial amount

f data obtained during the past two decades. It begins withn overview of the existing bulk composition models and theirpproaches, summarizes the datasets available and their util-ty, discusses the likelihood that data from the surface provideslobal information, and provides a complete reassessment of theulk composition of Mars in light of new data. I emphasize theomposition of bulk silicate Mars, but briefly discuss modelsor core composition, which are geochemically less well con-trained.

. Approaches to estimating bulk composition

Three main approaches have been used to estimate the bulkhemical composition of Mars. One takes a cosmochemicalpproach by defining different components and determining theirbundances in Mars. Another uses isotopic compositions of Mar-

ian meteorites and chondrite groups to define the abundancesf the chondritic raw materials that accreted to Mars. A thirdses broad geophysical properties to define mantle mineralogy.ote that these approaches have tended to be dominated by

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 418

geochemistry, isotopic geochemistry, or geophysics, but there isno reason why a blended model cannot be used.

2.1. Models based on geochemistry and nebular components

2.1.1. Wänke and Dreibus modelWänke (1981, 1987), Dreibus and Wänke (1984, 1987), Wänke

and Dreibus (1988, 1994), Longhi et al. (1992), and Halliday et al.(2001) estimate the bulk composition from element correlationsin Martian meteorites, with the assumption that refractory ele-ments are present in chondritic abundances. This model is directlytied to Mars through element abundances in Martian meteorites,and this direct link to samples of Mars has made the modelexceedingly robust, explaining why it is generally considered tobe the standard model for Martian bulk composition. Wänke andDreibus assume all elements more refractory than Mn are in chon-dritic proportions. From the cosmochemical viewpoint, refractoryand volatile tendencies are related to their condensation tem-perature in the solar nebula. The rough definitions of refractoryand volatile and some gradations are shown in Table 1, using50% condensation temperatures from Lodders (2003), calculatedat a pressure of 10−4 atm. The 50% condensation temperatureis simply the temperature at which half the mass of an ele-ment has condensed (the remainder is in the gas). The standardWänke–Dreibus model for Mars bulk composition is shown inTables 2 and 3.

As an example of the Wänke–Dreibus approach, consider howthey derive the FeO content of the mantle from the FeO/MnO ratioin meteorites. These oxides do not fractionate from each othersignificantly if the major phases are olivine and pyroxene (partitioncoefficients are about 1 for both), so igneous processes tend topreserve their mantle values. MnO values in Martian meteoritesare in the range 0.4–0.6 wt%. The CI chondrite MnO concentrationis 0.46 wt%, implying that Mars is not depleted in MnO. FeO/MnO inMartian meteorites is 39.1 (but see in Section 5), and FeO/MnO inCI chondrites is 100.6. Thus, if MnO is at CI abundance (0.46), thenFeO in the mantle is 39.1 × 0.46 = 17.9 wt%, in reasonable agree-ment with bulk Martian meteorite and GRS orbital data (Tayloret al., 2006a), and with models derived from the moment of inertia

(Bertka and Fei, 1998a,b). Other correlations include, for example,volatile elements with refractory elements (K/La, K/Th), correla-tions among the alkalis (K/Rb, Rb/Cs), and among the halogens(Br/Cl). Element correlations tend to be reliable for element pairs

G.J. Taylor / Chemie der Erde 73 (2013) 401– 420 403

Table 1Elements classified by geochemical behaviora and 50% condensations temperatures at a pressure of 10−4 atm (Lodders, 2003).

Category Temperature range (K) Lithophiles Siderophiles Chalcophiles

Refractory ≥∼1300 Zr, Hf, Sc, Y, lanthanides, Th,U, Al, Ti, Ta, Nb, Ca, Sr, Ba, V,Mg, Si, Cr, (Fe)

Re, Os, W, Ir, Mo, Ru, Pt,Rh, Ni, Co, (Fe), Pd

None

Moderately volatile 1230–800 P, Mn, K, Ga, Na, Clb, Rb, Cs, Au As, Cu, Ag, Sb, GeHighly volatile 750–250 F, Clb, Br, I None Bi, Pb, Zn, Te, Sn, Se, S, Cd, In, Tl, HgUltra volatile <182 H/H2O, N, C None None

a The classification into lithophile, siderophile, and chalcophile elements is not clear-cut. In the absence of metallic iron or sulfides, for example, elements behave likelithophile elements.

b Cl condensation temperature might be much lower than calculated by Lodders (2003), similar to those of Br and I.

Table 2Major element concentrations in models for bulk silicate Mars.

Wänke and Dreibusa Morgan and Andersb Sanloup et al.c Lodders and Fegleyd Khan and Connollye

SiO2 44.4 41.6 47.5 45.39 44TiO2 0.14 0.33 0.1 0.14 –Al2O 3.02 6.39 2.5 2.89 2.5Cr2O3 0.76 0.65 0.7 0.68 –FeO 17.9 15.85 17.7 17.21 17MnO 0.46 0.15 0.4 0.37 –MgO 30.2 29.78 27.3 29.71 33CaO 2.45 5.16 2.0 2.36 2.2Na2O 0.50 0.1 1.2 0.98 –K2O 0.04 0.01 – 0.11 –P2O3 0.16 – – 0.18 –

Total 100.03 100.02 99.4 100.00 98.7

a Wänke and Dreibus (1994).b Morgan and Anders (1979).

we

ce

TB

c Sanloup et al. (1999) model EH45:H55.d Lodders and Fegley (1997).e Khan and Connolly (2008).

ith similar bulk partition coefficients, such as the incompatible

lements K and Th or La, or the compatible elements Mg and Cr.

Inherent in the Wänke–Dreibus model is the concept that twoomponents combined to form Mars (and the other terrestrial plan-ts). It stems from ideas developed earlier by Ringwood (1979) and

able 3ulk composition of Martian crust + mantle (primitive Martian mantle).

W-Da M-Ab L-Fc

Be (ppb) 52 109 –

F (ppm) 32 23.6 41

Na (%) 0.37 0.071 0.73

Mg (%) 18.2 18.0 17.8

Al (%) 1.60 3.37 1.53

Si (%) 20.6 19.4 21.2

P (ppm) 700 1985 740

Cl (ppm) 38 0.88 150

K (ppm) 305 76.5 920

Ca (%) 1.75 3.68 1.68

Sc (ppm) 11.3 23.5 10.5

Ti (ppm) 840 1951 815

Cr (ppm) 5200 4469 4640

Mn (%) 0.36 0.116 0.284

Fe (%) 13.9 12.3 13.4

Co (ppm) 68 – 67

Ni (ppm) 400 – 140

Cu (ppm) 5.5 – 2.0

Zn (ppm) 62 41.9 83

Ga (ppm) 6.6 2.43 4.4

Br (ppb) 145 4.73 940

Rb (ppm) 1.06 0.26 3.5

Sr (ppm) 15.6 35.1 13.5

Y (ppm) 2.7 6.41 2.8

Zr (ppm) 7.2 38.0 8.3

Nb (ppb) 490 1938 –

a Wänke and Dreibus (1994), slightly embellished by Taylor and McLennan (2009).b Morgan and Anders (1979); siderophile and chalcophile elements not included.c Lodders and Fegley (1997); the siderophile/chalcophile elements P, Co, Ni, Cu, Ga, Mo

Wänke (1981). Component A is reduced and does not contain ele-

ments more volatile than Mn (this was originally formulated byWänke and Dreibus as more volatile than Na, but subsequent papersused the slightly more refractory element Mn as the cut-off). Themore refractory elements are in CI abundances. Fe and siderophile

W-Da M-Ab L-Fc

Mo (ppb) 118 – 17In (ppb) 14 0.10 12I (ppb) 32 0.59 120Cs (ppb) 70 25.9 154Ba (ppm) 4.5 9.88 5.4La (ppb) 480 926 400Ce (ppb) 1250 2457 1120Pr (ppb) 180 309 167Nd (ppb) 930 1704 850Sm (ppb) 300 506 250Eu (ppb) 114 194 99Gd (ppb) 400 691 395Tb (ppb) 76 130 69Dy (ppb) 500 877 454Ho (ppb) 110 193 98Er (ppb) 325 562 300Tm (ppb) 47 84 50Yb (ppb) 325 557 277Lu (ppb) 50 94 44Hf (ppb) 230 557 229Ta (ppb) 34 56 29W (ppb) 105 440 80Tl (ppb) 3.6 0.17 10Th (ppb) 56 125 56U (ppb) 16 35 16

, In, are Tl were calculated from metal-silicate partition coefficients.

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lements are metallic in Component A. Thus, the central assump-ion in this model is that the refractory elements are in CI relativebundances. Component B is oxidized and contains all elementsn CI chondritic abundances. These are useful constructs in under-tanding the components that accreted to the planets, although its not certain that the planets were really constructed from knownhondrites or their components. Morgan and Anders (1979) devel-ped a more elaborate multi-component model.

.1.2. Morgan and Anders modelGanapathy and Anders (1974) and Morgan and Anders (1979)

odeled Mars as a mixture of chondritic materials that hadeen modified by the same limited set of processes that affectedhondrites, such as variations in condensation temperature andractionation of metal from silicate. The model focuses on chon-ritic components formed in the solar nebula. This may not beorrect: as discussed below, Warren (2011) shows that terrestriallanets, differentiated meteorites, and non-carbonaceous chon-rites are clearly distinguishable from carbonaceous chondritesthe low-temperature component in the Morgan and Anders, 1979,

odel).Morgan and Anders (1979) proposed that there are three pri-

ary condensates from the solar nebula: a high-temperature,efractory-rich condensate; Fe–Ni metal; and magnesian silicates.organ and Anders (1979) define a fourth component, FeS and

eO, which they postulate formed by reaction of Fe metal with2S and H2O, respectively, with the FeO ending up in the magne-

ian silicate component. They also suggest a “remelted” component,y which they mean chondrules, hence depleted in volatiles (pre-umably lost during high-temperature chondrule formation). Theiragnesian silicates and chondrule components are not dramati-

ally different in composition. They further define a componentich in highly volatile elements, referred to in Morgan and Anders1980) as the “unremelted” component. When implemented tostimate the planet’s bulk composition, they use four main compo-ents, the refractories, magnesian silicates, metallic Fe, the volatileomponent (including both moderately volatile and highly volatilelements).

A central assumption in the Morgan and Anders approach is thatlements of similar volatility do not fractionate during nebular pro-esses, allowing them to use four “index” elements (U for refractorylements, Fe for metal, K for moderately volatile elements, and Tl or6Ar for highly volatile elements) to calculate the abundances of 83lements in the planet. When Morgan and Anders reported theirork in 1979, the idea of a group of meteorites being from Marsas just blossoming and quite controversial, so they did not useata from the meteorites. Their estimated Mars bulk composition

s given in Table 2.Morgan and Anders (1979) defined the concentrations of the

ndex elements for refractory (U, Th) and moderately volatileK) elements in Mars, using gamma-ray data from the Sovietrbiter Mars 5 and thermal models available at the time to predict

value of 620 for K/Th in bulk Mars, much lower than mea-ured by the Mars Odyssey gamma-ray spectrometer, GRS (5300;aylor et al., 2006a) or in Martian meteorites. In spite of theirstimate for K/Th being far too low, the basic approach is inter-sting and it is not significantly different from the Wänke–Dreibusethod. The idea that refractory elements are present at CI

hondrite relative abundances is common to both compositionalodels. The central difference is that Wänke and Dreibus, and my

eassessment of the Mars bulk composition (Section 5), use mul-iple elements determined independently, rather than assuming,or example, that the highly volatile elements, though depleted,ave CI (Wänke–Dreibus) or CV3 (Morgan and Anders) relativebundances.

e 73 (2013) 401– 420

2.2. Estimates based on isotopic composition of Martianmeteorites

Lodders and Fegley (1997), Sanloup et al. (1999), and Burbineand O’Brien (2004) focused on fitting the oxygen isotopic com-position of Mars, known from Martian meteorites, to mixturesof chondritic meteorites; their results are shown in Table 2 (andTable 3 for Lodders and Fegley’s model). Sanloup et al. (1999) pointout two important features of the isotopic approach. One is thatoxygen is the most abundant element and other models do notdetermine it explicitly. The other is that a model based on oxy-gen isotopes has only one major assumption, in contrast to severalwhen considering assorted components.

An isotopic estimate can also be tested by data from Mar-tian meteorites, orbiters, and landers. Lodders and Fegley’s (1997)assessment led to the estimate that Mars was constructed froma mixture of 85% H-chondrites, 11% CV-chondrites, and 4% CI-chondrites. In turn, this led to a predicted value of 16,000 for K/Thin bulk Mars, much higher than GRS data indicate (5300; Tayloret al., 2006a). Sanloup et al. (1999) took a similar approach inestimating the composition of Mars, arriving at a best fit beinga mixture of 45% EH and 55% H chondrites. They did not esti-mate the abundances of K and Th, but judging from their relativelyhigh estimated Na content, Sanloup et al. (1999) composition,appears to be enriched in moderately volatile elements comparedto Mars.

Burbine and O’Brien (2004) ambitiously examined over 225 mil-lion combinations of oxygen isotopic and chemical compositions of13 chondrite groups, testing reasonableness by comparison withassumed bulk FeO and Mg/Si and Al/Si for Mars (based on Mar-tian meteorites). They extracted only limited major element dataas their main goal was to test the feasibility that known mete-orite types could be mixed to produce Earth and Mars. Burbineand O’Brien (2004) find that the average of reasonable fits to Marsbulk composition could involve contributions from all 13 chon-drite types modeled, but are dominated by enstatite and ordinarychondrites, which make up about 60% of the contributors (carbona-ceous chondrites make up 25% and R chondrites contribute 14%).One problem with this approach is that Mg/Si and Al/Si are notreliable parameters to distinguish Earth and Mars (Filiberto et al.,2006; McSween et al., 2009). Martian rocks, for example, rangefrom distinctly lower than Earth in Al/Si in Martian meteorites toEarthlike or higher in rocks and soils at the Pathfinder and MERlanding sites and in global GRS data (see Fig. 3 in McSween et al.,2009).

Warren (2011) took a unique approach by using stable isotopesto evaluate mixing within the early solar system, specifically theisotopes of Ti, Cr, and O. His results support the notion that Mars(and Earth) could be mixtures of non-carbonaceous meteorites. Henotes that for Cr and Ti especially there is a dichotomy betweenmaterials formed in the outer solar system (mainly carbonaceouschondrites) and those formed in the inner solar system (ordinarychondrites, differentiated meteorites, and the Earth, Moon, andMars). This important observation is consistent with the resultsof Sanloup et al. (1999) model, which suggests that Mars is madeexclusively of non-carbonaceous meteorites, and with the solutionsdiscovered by Burbine and O’Brien (2004), which involve only 25%carbonaceous chondrites. Lodders and Fegley’s (1997) best esti-mate is also consistent as it indicates that the primary ingredientsof Mars involved 85% H chondrites and 15% carbonaceous chon-drites; the percentage of the carbonaceous chondrites componentis within the limits Warren (2011) estimates for the carbonaceous

contribution to Mars. As noted, the major problem with isotopicmodels for Martian bulk composition is that while isotopic com-positions can be matched, the abundances of volatile elements areover estimated.

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.3. Estimates based on geophysical properties of Mars

Geophysical properties such as the mean density, moment ofnertia, tidal Love number and dissipation factor, and radius providendependent information about the bulk composition of Mars,ncluding the size and composition of the core. We do not yet haveeismic data or heat flow for Mars, two important measurementshat will greatly enhance our knowledge of the Martian interior.owever, the InSight (Interior Exploration using Seismic Investiga-

ions, Geodesy and Heat Transport) mission, to be launched in 2016,ill make both seismic and heat flow measurements. Even without

hese data, geophysical studies have contributed to an improvednderstanding of the Martian interior and its composition.

One geophysical approach is to begin with a compositionalodel (usually the Wänke–Dreibus model) and calculate (e.g.,cGetchin and Smyth, 1978; Longhi et al., 1992; Sohl and Spohn,

997; Sanloup et al., 1999) or do experiments (e.g., Bertka and Fei,997) to determine how the mineralogy varies with pressure andemperature inside the planet. The results can be used to calculateulk density and moment of inertia for comparison with Martianeophysical properties.

A second geophysical approach is to invert geophysical datao construct mineralogical models of the interior, including phasehanges (Khan and Connolly, 2008). This approach has the advan-age of determining the bulk chemical composition for majorlements from mineralogy and mineralogical variation with depthpressure), with few assumptions. It does not assume CI chondritebundances, a particular mix of chondrites, or that we can inferomposition from geochemical correlations. It does not, of course,ive us minor and trace element concentrations. Nevertheless, it isf great utility and serves as an independent monitor of geochem-cal calculations.

Khan and Connolly (2008) use Gibbs energy minimization inhe system CaO–FeO–MgO–Al2O3–SiO2 (which make up 98% ofulk silicate Mars) to calculate the equilibrium mineralogy as aunction of depth and temperature inside Mars. The minimizationechnique to compute phase equilibria are described in detail byonnolly (2005). The results of the calculations, constrained byeasured geophysical properties (mean density, moment of iner-

ia, Love number, tidal dissipation factor) are shown in Table 2.he major element oxide concentrations are quite similar to theänke–Dreibus composition, though MgO is distinctly higher in

he composition determined by Khan and Connolly (2008), givingg/Si higher than in CI chondrites, which have the highest Mg/Si

f any chondrite group.

.4. Summary of the models

The Wänke–Dreibus geochemical model has become thetandard for Mars bulk chemical composition. The approach usedy Morgan and Anders (1979) is not fundamentally different. Itsstimate is quite different from that given by Wänke and Dreibus,articularly for volatile elements for which Morgan and Anders1979) values are significantly lower. As noted above, for example,

organ and Anders (1979) estimate K/Th of 620, considerably lesshan the global surface value measured by the Mars Odyssey GRS,300 (Taylor et al., 2006a). However, if good GRS measurementsnd Martian meteorites had been available to Morgan and Anders1979), most of the differences between their estimate and that of

änke and Dreibus would be minimal. The essential point is thatoth geochemical models depict Mars as composed of a refractory

omponent (condensation temperatures equal to or higher thanhat of Mn) present in chondritic (specifically CI) relative abun-ance, Fe partitioned between FeO and metallic Fe, and volatilesepleted.

e 73 (2013) 401– 420 405

The oxygen isotopic approaches have the great virtue of havingonly one parameter to match among mixtures of chondrite groups,and oxygen is the most abundant element. Major element oxidecompositions do not differ much from those of the geochemicalmodels, though the fact is that chondrites, like planetary mantles,are all ultramafic rocks. The significant difference between isotope-based and element-based estimates is the strong enrichment involatile elements in the isotope models (Tables 2 and 3). Thisenrichment is not seen in the GRS data: K/Th is 5300 for the Martiansurface versus 16,400 in Lodders and Fegley’s (1997) model. Sim-ilar differences are seen in the estimated abundances of the otheralkalis, halogens, and highly volatile elements such as Tl. Loddersand Fegley (1997) suggest that aqueous leaching in the mantle andhydrothermal alteration in the crust redistributed the volatile ele-ments. Taylor et al. (2006a,b) argue against this concept on the basisof only modest variations in K/Th in the crust. It is possible that thematerials that mixed to make Mars were like chondritic meteorites,as Lodders and Fegley (1997) and Sanloup et al. (1999) propose, butthat these components had not yet acquired their full complementof volatiles. Alternatively, volatiles could have been lost during theaccretion process (e.g., O’Neill and Palme, 2008). More likely, theaccreting protoplanets were differentiated, hence with the charac-teristics of differentiated meteorites (e.g., lower volatile contentsthan chondrites).

It is reassuring that the estimates based on geophysical proper-ties of Mars are similar to those obtained by the other independentmethods (Table 2). Although MgO is somewhat higher, it is prob-ably within error of the other estimates, and FeO (17 wt%) falls inthe narrow range of 15.8–17.9 wt%. The geophysical data almostcertainly provide the most reasonable estimate for core size andcomposition, as discussed in Section 6.

3. Datasets

Since the groundbreaking work of Wänke and Dreibus, substan-tial amounts of data have become available to allow a reassessmentof the composition of Mars from elemental data. Instead of only ∼10Martian meteorites that informed Wänke and Dreibus’ work, wenow have over 60 distinct Martian meteorites. Chemical, petrologic,mineralogic, and isotopic data are compiled in the Martian MeteoriteCompendium (Meyer, 2013), a crucially valuable resource. I haveused the data reported in the Compendium to compile a database ofMartian meteorite compositions. For each meteorite, the databaseconsists of averages of multiple analyses of the same meteorite,but for certain sets of elements, such as many volatile chalcophileelements, data from only one study are available. A particularlyuseful paper besides the Compendium is the careful summary ofcompositional data available through 1997 provided by Lodders(1998).

The meteorite data have the highest fidelity of any data avail-able, but are restricted to the random collection of Martianmeteorites. There is a bias toward younger ages in the collection,with all shergottites having ages < 500 Ga and the nakhlites andChasigny have ages of only 1.3 Ga. The only older ages reportedthus far are 4.1 Ga for ALH 84001 (Terada et al., 2003; Bouvieret al., 2009; Lapen et al., 2010) and 2.1 Ga for NWA 7034 (Ageeet al., 2013). Additional data come from orbiting spacecraft androvers. Of particular importance for understanding Martian bulkchemical properties are the data from the Mars Odyssey gamma-ray spectrometer (GRS). This instrument operated for eight yearsin Martian orbit, resulting in global data for K and Th, and equa-

torial data (from approximately 50◦ N to 50◦ S latitude) for Fe,Si, and H, with less quantitative information for S, Ca, and Al.Nominal spatial resolution is 5◦, about 500 km. While not impres-sive spatial resolution by the standards of imaging spectrometers,

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t is ideal for global data analysis problems. The GRS probeso a depth of a few tens of centimeters. GRS data are sum-

arized in Boynton et al. (2007, 2008) and available online atttp://geo.pds.nasa.gov/missions/odyssey/grs.html.

The highly successful Pathfinder and Mars Exploration Roverissions provide excellent data on rocks on the surface. Elemental

ata were obtained by the alpha particle X-ray spectrome-er (APXS). The rocks abraded with the rock abrasion toolRAT) on the MER missions are particularly useful as theyrovide complementary data to those provided by meteoritesnd the abrasion process limits that amount of soil maskinghe underlying rock. Soil analyses by themselves are useful,specially for assessing global average elemental abundances.oil data are available online at http://pds-geosciences.wustl.edu/issions/mer/geo mer datasets.htm. A good compilation and dis-

ussion of the data can be found in Brückner et al. (2008).

. Complications

.1. Is the Martian surface composition representative of thentire crust?

A significant uncertainty about estimating the compositions ofhe crust and mantle source regions is the extent to which wean use data obtained primarily from surface samples (meteorites,ocks and soils at the landing sites, orbital data) to determine theomposition of the entire crust and mantle. Taylor et al. (2006a)rgue that we can use the surface value for certain parameters,uch as the K/Th ratio, to represent the entire crust (and to a greatxtent the mantle). I summarize these and additional argumentsriefly here.

.1.1. Highlands megaregolithA substantial fraction of the crust (possibly about half;

cLennan and Grotzinger, 2008; Taylor and McLennan, 2009) wasonstructed before the end of the heavy bombardment at ∼3.8 Ga.t would have been repeatedly excavated and mixed by impacts,specially by basin-forming events, resulting in a thick (10–20 km)egaregolith. This early period of regolith formation may have

rovided a substantial fraction of the present-day regolith in thebsence of significant crustal recycling on Mars; that is, what goesnto the Martian crust stays in the Martian crust. Thus, soils andrbital (GRS) data may be sampling more than just the most recenturface rocks and sediments.

.1.2. Lava flowsLava flows visible at the surface of Mars are undoubtedly accom-

anied by larger volumes of intrusive rock stalled inside the crust.n Earth, the ratio of the intrusive to extrusive magma volumes is:1 (Crisp, 1984) in oceanic (basaltic) regions. If this holds for Mars,he abundant lava flows visible at the surface (themselves forminghick sequences of lavas) are accompanied by five times as muchntrusive magmas with similar composition. Furthermore, magmaompositions reflect the compositions of their mantle sources,ssuming we can correct for fractionation processes as the magmasigrated to the surface or stalled in magma chambers. Of course,e see only the uppermost, youngest lava flows, and it is possible

hat magma compositions changed with time. In fact, on the basis ofars Odyssey GRS data, Hahn et al. (2007) suggested that changes

n surface compositions reflect changes in magmatic compositions

ith age. Nevertheless, the compositions of all available lava flows

meteorites, Gusev rocks, inferred from GRS) are highly informativebout the compositions of the magmas from which the crust wasonstructed.

e 73 (2013) 401– 420

4.1.3. Element ratiosThe ratios of elements with very similar geochemical behavior in

igneous systems will reflect their ratio in the mantle. For example,K and Th do not readily fractionate, as shown by their similar, andvery low (�1), crystal-melt partition coefficients (Beattie, 1993;Borg and Draper, 2003; Hauri et al., 1994). Both elements are incom-patible and their concentrations in magmas are not greatly affectedby source rock composition or crystallizing phase, even when gar-net is involved. There are interesting exceptions, however. Th ishighly compatible in phosphate minerals (Jones, 1995). Phosphatesform late in the crystallization of a magma and are unlikely to beretained in a mantle source region, so probably do not play a role infractioning K from Th during igneous processes. However, in prin-ciple, it could be significant if mantle regions were metasomatizedby fluids that contained phosphate components. K is compatiblein phlogopite (Halliday et al., 1995) and somewhat compatible inamphibole (Halliday et al., 1995), so if these phases were presentin the Martian mantle, it could lead to fractionation of K from Th.In addition, a rock rich in K (named Jake M) has been analyzed inGale Crater, and classified as a mugearite by Stolper et al. (2013).Terrestrial mugerites form by extensive fractional crystallization ofalkaline basalts and are not abundant on Earth, implying that theirformation, if accompanied by fractionation of K from Th, are notof global compositional importance. No Th data are available forthe rock. Nevertheless, in general, K/Th in a lava flow reflects theratio in its mantle source region. Similar arguments can be madefor other element pairs (see Section 5). Studies of terrestrial basaltsshow that certain elements correlate strongly with one another(e.g., Jenner and O’Neill, 2012), implying that they have the samebulk partition coefficient. Following Wänke and Dreibus, I use suchcorrelations as guides to searching for correlations in the Martianmeteorite dataset.

4.2. Heterogeneity of the Martian mantle

The shergottites exhibit a large range in geochemical andisotopic composition. Their rare earth element (REE) patternsrange from severely light REE-depleted (CI-normalized La/Yb ∼0.1)with low abundances through to very slightly LREE-enriched (CI-normalized La/Yb ∼1.2) with high abundances. The variations inREE patterns correlate with isotopic compositions such as ini-tial 87Sr/86Sr and 143Nd/144Nd (e.g., Norman, 1999, 2002) andgeochemical parameters such as ratios among incompatible ele-ments (e.g., McLennan, 2003) and oxygen fugacity (e.g., Wadhwa,2001; Herd et al., 2002). These variations have been interpretedto indicate that the shergottites represent a mixture of two dis-tinctive sources. Nakhlites are enriched in incompatible elements(CI-normalized La/Yb ∼3) and isotopic data suggest a distinctivesource region (Foley et al., 2005). The origin of these distinctivesources is important. A discussion of how they might have formedis beyond the scope of this paper, but just their existence is impor-tant for understanding how to extract the bulk composition of Marsfrom element ratios.

In addition, Martian meteorites may not be a representativesample of the Martian crust or representative probes of its man-tle. This is made particularly clear in Fig. 1, a plot of the K/Th versusK. Virtually all the meteorites are lower in K and in K/Th than theGRS global mean. This suggests that we have not sampled a signifi-cant mantle source region that is richer in volatile elements than thesources represented by the meteorites. The two prominent samples(Nakhla and lherzolite NWA 1950) with K/Th around 10,000 may

be anomalous or affected by cumulate processes that fractionatedK from Th. Their high K/Th does not appear to be due to terres-trial weathering: Nakhla is a fall and both have Th/U > 2, whereassamples altered on Earth have Th/U < 0.2. The central point is that

G.J. Taylor / Chemie der Erd

Fig. 1. Trace element characteristics of Martian meteorites and the average bulksurface (Taylor et al., 2006a) suggest diverse sources in the Martian mantle. Por-tions of the mantle must have higher K/Th to counter the lower values in meteoritescompared to the Mars Odyssey gamma-ray spectrometer (GRS) mean surface com-p2t

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osition. Similar characteristics are shown by other trace elements (McLennan,003; Taylor et al., 2008). This mantle diversity adds an element of uncertaintyo determinations of bulk composition from trace element ratios.

ars may contain a reservoir with higher K/Th than the GRS globalean, needed to balance the low K/Th of the meteorites (Brandon

t al., 2012). The high K/Th ratio of the Martian surface measured byRS might be caused by secondary aqueous alteration processes,

mplying that this ratio is not reflective of the bulk crust. Taylort al. (2006b) explored this possibility and although plausible, apecific mechanism giving rise to a planetary-scale change in K/Thatio of surficial materials was not identified. The existence of these

able 4ulk composition in revised model for bulk silicate Mars.

Conc 2-� Methoda Ratiob

Li (ppm) 3.0 1.7 S-B Li/Fe

Be (ppb) 47.7 0.4 S-B Be/Nd

F (ppm) 21 13 A-B F/K

Na (%) 0.40 0.08 S-A Na/Al

Mg (%) 18.5 0.7 R –

Al (%) 1.64 0.04 R –

Si (%) 20.5 0.9 R –

P (ppm) 675 215 S-A P/Yb

Cl (ppm) 32 9 S-A Cl/Th

K (ppm) 309 36 A(GRS)

Ca (%) 1.74 0.04 R –

Sc (ppm) 11.0 0.4 R –

Ti (ppm) 832 30 R –

Cr (ppm) 4990 420 R –

Mn (%) 0.34 0.05 S-A, D –

Fe (%) 14.1 0.8 S-A Fe/Mn

Co (ppm) 71 25 S-A Co/Ni

Ni (ppm) 330 109 S-A Ni/Mg

Cu (ppm) 2.0 0.7 S-A Cu/Mg

Zn (ppm) 18.9 2.9 A-A Zn/Sc

Ga (ppm) 6.6 0.8 S-A Ga/Al

As (ppb) 86 55 A-B As/Ce

Se (ppb) 85 36 S-A Se/Yb

Br (ppb) 191 58 S-A Br/Cl

Rb (ppm) 1.30 0.14 S-S Rb/La

Sr (ppm) 14.6 0.7 R –

Y (ppm) 2.89 0.52 R –

Zr (ppm) 7.49 0.60 R –

Nb (ppb) 501 0.07 R –

Ru (ppb) 2.6 0.9 A-Ac Ru/Mg

Pd (ppb) 2.4 0.8 A-Ac Pd/Mg

Ag (ppb) 4.2 2.8 A-A Ag/Dy

a R: refractory element from volatile-free CI chondrite composition. S-B: slope of correor shergottites, including lherzolitic shergottites. S-A: slope of correlation line for all mlivine-phyric and basaltic shergottites. A-A: average of ratio to abundance of another elef surface. D: partition coefficient for Mn in basaltic melt divided by Mn in peridotite, usib Ratio of element to refractory element used in slope or average methods.c Based on concentrations of element in samples containing ≥15 wt% MgO (Brandon et

e 73 (2013) 401– 420 407

reservoirs indicates that some caution is advised when extract-ing planetary bulk compositions from a non-representative set ofsamples. This problem is further complicated by the role of vari-able oxidation state throughout the mantle, which might explainsome differences between the Martian meteorites and the Gusevrocks (Tuff et al., 2013), but also changes geochemical behavior ofsome trace elements, making the link from meteorite to mantlecomposition less clear.

5. A reassessment of Martian bulk composition

I take a geochemical approach to reassess the bulk composi-tion of Mars using the methods outlined in Longhi et al. (1992),with emphasis on its bulk silicate composition. It is basically thesame approach as reported in the insightful papers by Wänkeand Dreibus, but with the addition of estimates of the uncertain-ties for each element and to show more of the correlations orlack of correlations than usually given. It also minimizes assump-tions about condensation temperatures (e.g., components A and B),except for the unavoidable assumption that refractory elements arepresent in CI relative abundances. The new (only slightly modified)chemically-determined bulk composition is given in Tables 4 and 5.

5.1. Estimating uncertainties

Estimates of the uncertainties in Martian bulk compositionstem from analytical and sampling uncertainties for each sam-ple or the global GRS data, and from variations from sample tosample. These can be quantified from the data. Uncertainties also

Conc 2-� Method Ratio

Cd (ppb) 9.6 6.1 A-A Cd/DyIn (ppb) 6.9 2.2 S-A In/YSn (ppb) 38.5 7.0 A-A Sn/SmI (ppb) 36 22 A-A I/ClCs (ppb) 95 37 S-S Cs/LaBa (ppm) 4.37 0.21 R –La (ppb) 439 48 R –Ce (ppb) 1170 110 R –Pr (ppb) 176 5.8 R –Nd (ppb) 864 60 R –Sm (ppb) 274 13 R –Eu (ppb) 103 6 R –Gd (ppb) 374 31 R –Tb (ppb) 67.3 13.2 R –Dy (ppb) 450 17 R –Ho (ppb) 106 6 R –Er (ppb) 306 37 R –Tm (ppb) 44.8 6.2 R –Yb (ppb) 308 26 R –Lu (ppb) 44.8 1.7 R –Hf (ppb) 217 14 R –Ta (ppb) 27.2 1 R –W (ppb) 74 31 S-A W/ThRe (ppb) 0.88 0.66 A-Ac Re/MgOs (ppb) 2.0 0.8 A-Ac Os/MgIr (ppb) 2.0 1.0 A-Ac Ir/MgPt (ppb) 3.1 0.8 A-Ac Pt/MgTl (ppb) 1.28 0.71 S-A Tl/ThBi (ppb) 0.60 0.41 A-A Bi/ThTh (ppb) 58 12 R –U (ppb) 16 3 R –

lation line for olivine-phyric and basaltic shergottites. S-S: slope of correlation lineartian meteorite types. A-B: average of ratio to abundance of another element forment for all martian meteorites. A(GRS): Average Mars Odyssey GRS K, Th analysisng experimental data.

al., 2012).

408 G.J. Taylor / Chemie der Erd

Table 5Major element concentrations in revised model for bulk silicate Mars.

Concentration (wt%) Uncertainty ±2 sigma Methoda Ratiob

SiO2 43.7 1.0 Refrac –TiO2 0.14 0.01 Refrac –Al2O 3.04 0.10 Refrac –Cr2O3 0.73 0.04 Refrac –FeO 18.1 1.0 Slope Fe/MnMnO 0.44 0.06 D –MgO 30.5 0.05 Refrac –CaO 2.43 0.01 Refrac –Na2O 0.53 0.10 Slope Na/AlK2O 0.04 0.002 Average K/ThP2O3 0.15 0.047 Slope P/YbTotal 99.8

a Refrac: refractory element from volatile-free CI chondrite composition. Slope:slope of correlation line. Average: average of ratio to abundance of another element.De

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: partition coefficient for Mn in basaltic melt divided by Mn in peridotite, usingxperimental data.b Ratio of element to refractory element used in slope or average methods.

rise from assumptions made in determining the bulk composi-ion, such as the assumption that elements more refractory than

n are present in CI relative abundances. I show below that tohe extent we can test this assumption, the relative abundancesf refractory elements are in chondritic proportions. Anotherssumption is that a good correlation between two elements iseaningful. This would seem to be strongest when the corre-

ations hold for datasets involving meteorites from all mantleources.

I assessed the analytical uncertainties in several ways: Forhe refractory elements, I simply used the 2-sigma uncertain-ies determined by Lodders (2003) for CI chondrite analyses. For

n I used the uncertainty in experimental determinations of theartition coefficients between basaltic melts and peridotites at aange of pressures. These were determined by fitting a line to thexperimental data, using the least squares method outlined byork (1969). (The least-squares calculations used a spreadsheetenerously provided by Randy Korotev, Washington Universityn St. Louis). The calculation fits a line and determines 2-sigmancertainty in its slope and intercept. Similarly, for FeO, I usedxperimental data and fit a line of the crystal/liquid partition coef-cient for FeO, and its 2-sigma uncertainty.

For trace elements I used linear correlations when the squaref the correlation coefficient (R2) was equal to or greater than.5, and applied the York approach to determine the uncertainty

n the slope. Elements usually involved a refractory element

ig. 2. Abundances of lithophile elements (Lodders and Fegley, 1998) in chondrite groupn increase for the two most refractory elements (Zr and Hf), relative abundances are esshat these elements are present in chondritic relative abundances in Mars.

e 73 (2013) 401– 420

(whose composition is known through the assumption of uniformrefractory element abundances) and a volatile one. In most cases,I was able to fit a line through the origin, at least to within theuncertainty of the intercept. That is, I forced the line through a zerointercept, allowing the slope to serve as a proxy for the averageratio of the two elements plotted. This approach assumes thatthe scatter about a line is caused by the combination of analyticalprecision, sampling errors, natural variation within the sample set,and small differences in bulk partition coefficients (Hanson, 1989).

In some cases, R2 was less than 0.5, but the data did notscatter across a large compositional space (typically a range ofless than a factor of two of the mean). In those clustered but notlinearly correlated cases I calculated the standard errors of themean of each element. The uncertainty in the ratio is given byRa/b [(�a/a)2 + (�b/b)2]1/2, where Ra/b is the ratio of two elementswith concentrations a and b, and �a and �b are the standard errorsfor elements a and b. In cases where a two-element comparisonwas simply too scattered to be meaningful, I did not determine theconcentration of the unknown element.

5.2. Refractory element abundances

The cosmochemical approach assumes that refractory elementsare present in Mars at chondritic relative abundances. Below, Iassume that all elements with higher 50% condensation temper-atures than that of Cr are present at CI relative abundances. Oneway of testing this crucial, and common, assumption is to exam-ine the abundances of the refractory lithophile elements in majorchondrite groups, normalized to CI abundances (Fig. 2). In spite ofsome irregularities in the abundance patterns, the general impres-sion is that the elements are largely unfractionated, as expected.Differences in the compositions of chondrite groups reflect differ-ences in the histories and relative abundances of their constituentsin the solar nebula and in their parent bodies (e.g., Huss et al., 2003).There is a curious increase in relative abundances from Y to themore refractory elements, with roughly similar slopes. (Note thatthe concentration axis in Fig. 2 is linear rather than the conven-tional logarithmic abundance.) There is also a slight decrease at theleast refractory end, from Eu to Cr. The slight depletion in Cr, Si, andMg is on the order of only 10%, which in principle could translateto a 10% uncertainty (or perhaps a systematic over estimate) in the

abundances of these elements in Mars.

I also tested element correlations among the refractory elementsin Martian meteorites (Fig. 3), using element pairs with similargeochemical behavior. The incompatible trace element pairs Th/La,

s normalized to CI chondrites (using the mean given by Lodders, 2003). Except forentially unfractionated. This gives some confidence that it is reasonable to assume

G.J. Taylor / Chemie der Erde 73 (2013) 401– 420 409

F t thatc repor

ZTrtc3tta

Foo(

ig. 3. Correlations among pairs of refractory elements, plotted to test the conceporrelations are within errors of the ratio in carbonaceous chondrites (using values

r/Hf, and Zr/Y have reasonably linear slopes with R2 > 0.7 (Fig. 3).he ratios are within the 2-sigma uncertainties of the CI ratios. Theatio of the compatible major element oxides Cr2O3/MgO forms aight linear array (R2 of 0.91) with a statistically significant inter-ept. When the line is extrapolated to the bulk planetary MgO of

0.5 (Table 5), the MgO/Cr2O3 ratio is 30.3 ± 4.9, within error ofhe CI value of 26.5. While not overwhelmingly convincing, theseests are consistent with the assumption that refractory elementsre present at CI relative abundances in Mars.

ig. 4. FeO versus MnO in Martian meteorites and rocks at the Gusev landing siten Mars (Ratted: abraded rock samples; Brushed: less abraded rock samples). Gusevlivine basalts and alkaline basalts are from a data analysis by McSween et al.2006a,b).

refractory elements in Mars are present in chondritic proportions. Slopes of theted by Lodders, 2003).

Using the cosmochemical approach, I calculated the concen-trations of all elements equal to or more refractory than Cr, andrenormalized the sum of those elements to 100. It excludes Ni andCo, which concentrate in metallic iron during core formation. Italso ignores oxygen at this stage. A complication arises for Fe as itis partitioned between both FeO and metallic Fe in Mars, indicatingthat in the calculation of refractory elements in bulk silicate Marswe should use only the Fe in silicate Mars. This is derived fromFeO in the crust (see next section). All elements are then convertedto oxides; oxide concentrations in bulk silicate Mars are given inTable 4. The uncertainties reflect the uncertainties in the averageCI chondrite value and the uncertainty in FeO (see below), whichmatters because of the normalization to 100%. Compared to Wänkeand Dreibus (1994), my revision is slightly lower in SiO2.

5.3. FeO and MnO

Wänke and Dreibus derived the FeO concentration in Mars fromthe MnO content in Martian meteorites. They noted that MnO haspartition coefficients for major minerals and melts close to 1, soits abundance in the mantle is the same as in the crust. The setof meteorites available to them had an average MnO content of0.46 wt%. Noting that this was the same as in CI chondrites, and thatFeO/MnO in CI chondrites was 100.6, they used the FeO/MnO ratioin Martian meteorites (39.5 ± 1.2) and the MnO in the meteoritesto derive a bulk FeO of 17.9 ± 0.6.

This approach is rigorous, but there are two problems. One isthat the partition coefficient (MnO in peridotite/basaltic melt) isnot exactly 1.0. Using experimental data reported by Takahashiand Kushiro (1983), Walter (1998), Hertzberg and Zhang (1996),

4 er Erde 73 (2013) 401– 420

WcaraVair(s

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Fig. 5. Histograms of FeO (wt%) distributions from the Mars Odyssey gamma-rayspectrometer (GRS, from 5-degree grid points, Taylor et al., 2006a) and (bottom)Martian meteorites, abraded (Ratted) rocks from the Gusev landing site, and meansoils from Gusev, Meridiani, Pathfinder, and Viking-1 landing sites. Note the peaks at

FeO concentration coupled with the wrong CaO/Al2O3 ratio. Theyconcluded that the solution was a polybaric differentiation process,

Fig. 6. Apparent partition coefficient (concentration in initial peridotite divided bycomposition in equilibrium melt) versus pressure in kilobars. Dashed lines brackettypical pressures thought to represent source depths for Martian magmas (Filibertoand Dasgupta, 2011). In this pressure range, the partition coefficient is close to unity,

10 G.J. Taylor / Chemie d

asylenki et al. (2003), and Le Roux et al. (2011), it is clear that therystal/liquid partition coefficient for MnO is less than 1.0. The aver-ge for measurements in the pressure range 10–25 kb (expectedange for magma generation), is 0.93 ± 0.08 (2-�). The revised aver-ge MnO in Martian meteorites is 0.48 ± 0.12 (ignoring Governadoraladares, which has an anomalously high MnO content in the onenalysis available). The partition coefficient and the average MnOn the crust implies a bulk planetary MnO of 0.44 ± 0.08. Using theevised FeO/MnO for the larger set of Martian meteorites, 40.3 ± 2.32-�), giving an FeO concentration of 17.7 ± 2.3 wt%, essentially theame as determined by Wänke and Dreibus (17.9 wt%)

In spite of the agreement between my new estimate using con-iderably more meteorites in the database, an important caution isn order. Igneous rocks analyzed at the Gusev landing site haveonsiderably higher FeO/MnO than that recorded by the mete-rites (Fig. 4), a point raised by McSween et al. (2009). Almost allnalyses of basaltic rocks at Gusev, whether brushed or abradedo remove adhering dust and weathering products, have highereO/MnO than do the meteorites. McSween et al. (2006a,b) deter-ined the compositions of two main suites of rocks at Gusev,

livine basalts and alkaline basalts, using all available data toxtrapolate measured compositions to the compositions of thelteration-free igneous rocks. These are also plotted in Fig. 4 andre slightly (olivine basalts) to significantly (alkaline basalts) offsetompared to Martian meteorites. This may reflect distinct sourceegions or oxidation state of the source regions (Tuff et al., 2013) forhe magmas from which the Gusev rocks formed, and raise ques-ions about the utility of using FeO/MnO to deduce bulk FeO in

ars.An alternative approach is to determine bulk FeO independent

f its correlation with MnO. All data indicate without question thathe Martian crust is richer in FeO than terrestrial basaltic rocks∼10 wt%) as shown by FeO concentrations measured by GRS, andn meteorites, soils, and abraded rocks (Fig. 5). FeO peaks at close to9 wt% with a global mean of 18.4 ± 0.2 (based on summing all spec-ra, with the 2-� uncertainty based on counting statistics (Taylort al., 2006a), but lower values, down to 13 wt% in some Gusevocks are significant. The values significantly below the meteoritend GRS distributions may indicate the presence of mantle regionsith lower FeO than typical.

Experimental data on FeO partitioning during melting of peri-otite provide a way to calculate the mantle FeO from the meanurface FeO. Fig. 6 shows data compiled from several studies (seeaption for references). It assumes that the solid/melt partitionoefficient for FeO is closely approximated by the ratio of FeO in theriginal peridotite solid to the FeO in the melt. Data used are only forases where the amount of melting is no more than 25%. Note therend of decreasing D with increasing pressure. In Mars, magmasppear to form at pressures of 10–25 kb (Filiberto and Dasgupta,011). In this range, the apparent D(FeO)S/L straddle the D = 1 line.sing the linear fit to the data, at 20 kb, D(FeO) is 0.95 ± 0.06. Using

he mean crustal FeO determined by the GRS, 18.4 wt%, and using(FeO) of 0.95, I calculate a mantle FeO of 18.1 ± 2.2 wt% (2�);

his is the value used in Table 5. This agrees within uncertaintiesith the result obtained from FeO/MnO of the Martian meteorites.ore importantly, the geochemical estimate agrees with the FeO

etermined by Khan and Connolly (2008) on the basis of bulk geo-hysical properties of the planet, 17 wt%. This agreement indicateshat lower FeO sources like those giving rise to some Gusev basalts

ake up a small fraction of the Martian mantle.Agee and Draper (2004) made an independent assessment of

he FeO concentration in the Martian interior through experi-

ents on an L-chondrite composition at 5 GPa pressure. Their

ntent was to try to determine if the mantle source rocks for thehergottite group of Martian meteorites could be formed from aelatively FeO-rich composition like L chondrites. The L-chondrite

18–19 wt%; even the lowest values are higher than typical Mid-Ocean Ridge basaltson Earth (∼10 wt%).

composition used was that of the Homestead chondrite, whichis similar to the Wänke–Dreibus bulk Mars composition. Theyfound that the Homestead composition produced magmas thatcontained too much FeO than even the FeO-rich shergottites con-tain, coupled with the correct CaO/Al2O3 ratio, or had a reasonable

implying that surface basalt FeO concentrations are reliable indicators of sourceregion (hence mantle) concentrations.

Data are from experimental studies by Falloon and Green (1988), Faloon et al.(1988), Takahashi and Kushiro (1983), Hertzberg and Zhang (1996), Walter (1998),Wasylenki et al. (2003), and Agee and Draper (2004).

G.J. Taylor / Chemie der Erde 73 (2013) 401– 420 411

Fig. 7. P versus refractory element Yb for olivine phyric and basaltic shergottites.Tvr

bH

rmc7mo2ooMsTMsatoMyt

5

trroPic

5

lsAc

ra

Fig. 8. Moderately volatile elements Ga and Na plotted against refractory element

he correlation coefficient indicates a significant correlation between moderatelyolatile P and refractory element Yb. Knowing Yb from the assumption of chondriticefractory abundance allows calculation of the P abundance in bulk Mars.

ut with a shergottites source containing an FeO content more like-chondrites than L-chondrites, ∼15 wt% FeO.

Concentrations (Table 5) of FeO (18.1 wt%) and MgO (30.5 wt%)esult in an Mg# [100 × molar Mg/(Mg + Fe)] of 75 ± 4. This is in theiddle of the range of Mg# in the most magnesian value in the

ores of olivines in olivine-phyric shergottites, which range from0 to 86 (Meyer, 2013). Thermal emission spectroscopy show thatagnesian olivine is exposed in places on Mars, the most notable

f which is a ring of the Argyre basin (Koeppen and Hamilton,008; Lane and Goodrich, 2010). Comparing to laboratory spectraf experimentally produced olivine suggests that the Argyre olivineccurrences are associated with olivine-rich basalts with olivineg# in the range 85 ± 5, within uncertainty of the most magne-

ian olivine in olivine-phyric basalts (86, Musselwhite et al., 2006).aken together, these data show that the Martian mantle varies ing#, but that the bulk composition is close to an Mg# of 75. The

omewhat lower FeO reported by Agee and Draper (2004) results in higher Mg#, ∼79. The geophysically-determined FeO concentra-ion of 17 wt% combined with the geochemically determined MgOf 30.5 produces a bulk Mg# of 76.8. Khan and Connolly’s (2008)gO is higher than the geochemically based one (Table 2), and

ields a bulk Mg# of 77.6. Even the highest of these values is lowerhan the terrestrial bulk Mg# of 89.

.4. Phosphorous

Phosphorous is a lithophile element with some siderophileendencies. It correlates well with Yb (Fig. 7). The slope of the cor-elation line gives the P/Yb ratio in Mars, assuming these elementseflect the bulk composition. Multiplying by the Martian bulk Ybf 308 ppb (Table 4) gives a P concentration of 675 ppm (0.15 wt%2O5). This is somewhat depleted compared to CI chondrites (P/CIs 0.7), perhaps indicating that P was partitioned partially into theore during primary Martian differentiation.

.5. Moderately volatile elements

Gallium and sodium correlate well with Al (Fig. 8) with corre-ation coefficients (R2) of 0.92 (Ga–Al) and 0.75 (Na–Al). Using thelopes, which provide Ga/Al and Na/Al, and the concentration ofl in Mars (Table 4, 1.63 wt%), I calculate that bulk silicate Mars

ontains 6.6 ppm Ga and 0.40 wt% Na (0.53 wt% Na2O).

Potassium is best determined in the GRS data as it more likelyepresents a global average. It correlates well with Th as bothre strongly incompatible large-ion lithophile elements, hence are

Al. Bulk Mars concentrations for Ga and Na can be determined from the Ga/Al andNa/Al ratios. Data represent analyses of all Martian meteorite types.

likely to reflect the bulk planetary ratio. K/Th varies across the Mar-tian surface (Fig. 9). Taylor et al. (2006b) examined the possiblereasons for this, with no definitive answers. It is clear, however,that the distribution is close to Gaussian and well defined. The meanK/Th is 5330 ± 440 (2-�). In this case, the uncertainty is calculatedfrom the sum of all the spectra obtained by the GRS, involving over2 × 107 s of counting time. Thus, it reflects the counting statistics,not the variation in the data, but does reflect the accuracy withwhich we know the mean. Using the K/Th ratio and the Th con-centration of 58 ppb (Table 4), I calculate a bulk K concentration of309 ± 26 ppm.

I obtain the concentrations of the other alkali elements by thecorrelations with La in shergottites only (Fig. 10A, C). Including theother Martian meteorites renders the correlations much weaker. Rbis particularly well correlated with La (Fig. 10A), revealing a Rb/Laratio of 2.91. Using a bulk La content of 439 ppb (Table 4) I esti-mate a bulk Rb concentration of 1.27 ± 0.13 ppm. Rb also stronglycorrelates with K and although the K/Rb ratio is commonly used toestimate planetary geochemical reservoirs, it is in principle betterto use the ratio to a refractory element (La) rather than to an ele-ment (K) that is itself determined by a ratio. Nevertheless, usingthe strong correlation between K and Rb (Fig. 10B), I estimate a Rbcontent of 1.45 ± 0.25, within error of the estimate using Rb/La. Thisgives some credence to restricting the Rb calculation to the olivinephyric and basaltic shergottites only. The Cs (Fig. 10C) data form agood linear array. Using the bulk La value (Table 4) and a slope of0.180, I estimate a Cs abundance of 79 ± 31 ppb.

Other moderately volatile elements include As, Cu, Ag, and Sb, allsomewhat chalcophile or siderophile. I found no significant corre-

lation between Sb and any element, and the plots were so scatteredthat an average of Sb with a geochemically similar element wouldnot be informative. It is not clear whether the lack of correlationis due to analytical issues or complicated geochemical behavior

412 G.J. Taylor / Chemie der Erde 73 (2013) 401– 420

F degreet are cac n and

iamA(cvc

5

tCdiamMdtBR

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ig. 9. Histogram from Mars Odyssey gamma-ray spectrometer determinations (5-he ratio in upper few tens of centimeters of surface. Global mean and uncertaintyounting time. Uncertainty represents the uncertainty in our knowledge of the mea

n Mars. As and Ag also do not form linear arrays when plottedgainst suitable elements, but do cluster sufficiently to allow esti-ating their abundance from averages, with Ce for As and Dy forg. Results are given in Table 4. Copper correlates well with Mg

Fig. 11), steadily declining with increasing MgO, indicating lessompatible behavior for Cu than Mg. I model the bulk Cu as thealue where MgO equals bulk Mars (30.5 wt%). This results in a Cuoncentration of 2.0 ± 0.7 ppm.

.6. Highly volatile elements, including halogens

The highly volatile elements have 50% condensation tempera-ures less than 750 K (Lodders, 2003), and include Bi, Zn, Sn, Se,d, In, and Tl, plus the halogens. (Other elements in this groupo not have good correlations with a refractory element. These

nclude Pb, Te, and Hg.) The two best-determined elements are Tlnd In (Fig. 12). Both correlate acceptably with a refractory ele-ent, from which we can determine their concentrations in bulkars (Tl is 1.4 ± 0.7 ppb, In is 6.9 ± 2.2 ppb). The other elements are

etermined from average values in Martian meteorites comparedo the average of an element with similar geochemical behavior:i/Th, Zn/Sc, Sn/Sm, Se/Yb (correlates with R2 of 0.5), and Cd/Dy.esults are given in Table 4.

The halogens are potentially valuable as they behave as incom-atible elements, hence their concentrations in the mantle can beetermined from their correlations with refractory incompatiblerace elements. Furthermore, Br and I have low condensation tem-eratures, thus potentially helping us determine the abundance ofll highly volatile elements. However, bulk Martian halogens areifficult to assess because, except for F, they are readily lost from

ava flow surfaces in gas phases, so our database of mostly extrusiveocks is ambiguous with regards to the magmatic source regions inhe mantle.

Because loss is so common for halogens from lavas, Taylor et al.2010) suggested that the mean surface Cl/K (1.27, determined byRS), which is close to the chondritic value of 1.28, may reflect

he bulk composition of silicate Mars. If the chondritic Cl/K is not coincidence (a distinct possibility considering that Cl is highlyobile in aqueous fluids and heterogeneously distributed on both

over and GRS scales), then the ratio of these two incompatible

lements suggests that the CI-normalized Cl abundance is about theame as that of K, 0.6. This is consistent with the concentration ofl on the surface, 0.5 wt%. Such an elevated Cl concentration is alsoonsistent with the similar condensation temperatures of Cl and K,

grid points) of K/Th ratio of Martian surface (Taylor et al., 2006a,b). Data representlculated on the basis of counting statistics of a global spectrum of over 2.4 × 107 s

does not reflect the natural variation in K/Th.

948 and 1006, respectively (Lodders, 2003). Furthermore, Cl and Brare well correlated at a ratio close to chondritic (Fig. 13B), implyingthat Br has a normalized abundance of 0.5, too. This is surprisingin light of low 50% condensation temperature of 546 K for Br. Thecondensation temperatures might be wrong, of course, as they aredependent of what phases are assumed to contain trace elements.Or the chondritic Cl/K is just a coincidence and Cl is actually lowerthan in bulk Mars than in the uppermost crust.

If we assume that the Martian meteorites did not lose their mag-matic Cl, then a plot of Th versus Cl (Fig. 13A) allows us to estimatea bulk Cl value of 32 ± 9 ppm. This is lower than Wánke and Dreibusdetermined, 38 ppm, but within the uncertainties. Given the goodcorrelation between Br and Cl (Fig. 13B), I estimate a Br concentra-tion of 191 ± 58 ppb. The Br/Cl ratio is close to chondritic (Fig. 13B).Iodine was determined from the average I/Cl in Martian meteorites(data are not well correlated). The normalized abundances of Cl, Br,and I are all in the range 0.05–0.07. Low halogens are supported byanalyses of melt inclusions in olivine phenocrysts in olivine-phyricshergottite Y 980459 (Usui et al., 2012). They report that only mod-est losses of F and Cl could have taken place because both correlatewell with Na in both melt inclusions and in the glassy groundmassof the rock. Thus, it seems safe to conclude that the Martian mete-orites in general record the halogen concentrations in the mantleand the original and Wánke and Dreibus approach is appropriate.

5.7. Ni and Co

Ni and Co correlate well with Mg (Fig. 14). From the ratio of eachto Mg and the Mg concentration in Mars (Table 4), we get a Ni con-centration of 330 ± 109 ppm and a Co concentration of 71 ± 25 ppm.These are lower than the bulk planet has (∼2 wt% and ∼0.1 wt%,respectively) because most of inventory of these elements is in thecore.

5.8. Strongly siderophile elements

Brandon et al. (2012) present high quality analyses of highlysiderophile elements (Os, Ir, Ru, Pt, and Re) in shergottites. Forthose with MgO greater than about 15 wt%, the siderophiles arein chondritic proportions. If those abundances represent the bulkcomposition of Mars, hence assuming that the siderophile elements

in the other shergottites with lower MgO have been fractionated,we can estimate the abundances in bulk silicate Mars from Brandonet al.’s (2012) data. I averaged the concentrations of the elements inshergottites measured by Brandon et al. (2012), and assumed that

G.J. Taylor / Chemie der Erde 73 (2013) 401– 420 413

Fig. 10. Correlation of moderately volatile alkali elements versus refractory La (A,C) and for Rb versus K (B). Correlations can be used to determine abundances of thealkali elements (see text). Data are for shergottites (including lherzolitic shergot-tites).

Fig. 11. Cu versus MgO concentrations correlate reasonably well, for all Martianmeteorite types. A Martian bulk silicate concentration for Cu can be calculated fromthe linear fit for Cu versus MgO at the point where MgO has the bulk silicate Marsvalue of 30.5 wt% (Table 5).

Fig. 12. Correlations of highly volatile elements versus refractory elements, for thefew data available for all Martian meteorite types. (A) Tl versus Th. (B) In versus Y.Correlations are only modestly strong, but sufficient to allow estimation of In andTl in bulk silicate Mars.

Fig. 13. Cl–Th (A) and Br–Cl (B) correlations for all Martian meteorites. Note thatBr/Cl is close to the chondritic ratio.

414 G.J. Taylor / Chemie der Erde 73 (2013) 401– 420

FuC

tcTt2neEt

5

tLoHMo5st

teWhsart9o

Fig. 15. H2O versus Cl for 5-degree grid points measured by Mars Odyssey gamma-ray spectrometer (GRS). GRS actually determines concentration of H, which has been

ig. 14. Siderophile elements Ni and Co versus Mg. Correlations are reasonable andseful for deriving Ni and Co concentrations in bulk silicate Mars. Most of the Ni ando inventories are probably in the core.

heir relative and absolute abundances reflect their mantle con-entrations. The results appear in Table 4 and are discussed below.his approach is probably sound for Pt because during silicate par-ial melting, Pt has a partition coefficient close to 1 (Jones et al.,003). All these elements have almost identical abundances whenormalized to CI chondrites (using the values for Orgueil, Horant al., 2003), ranging from 0.0036 to 0.0045, except for Re (0.22).xcluding Re, the group of elements averages ∼50% of the value inhe terrestrial primitive upper mantle (Becker et al., 2006).

.9. H2O and D/H

Observations of the Martian surface from spacecraft revealedhat its crust has been modified by substantial fluvial activity.anded spacecraft and studies of weathering veins in Martian mete-rites have repeatedly confirmed the presence of water in Mars.owever, Dreibus and Wänke (1987) estimate that bulk silicatears contains only 39 ppm of H2O, the equivalent of about 15%

f a terrestrial ocean-equivalent and 7% of the lower estimate of00 ppm in the bulk Earth (see review by Mottl et al., 2007). Thiseems insufficient considering the evidence for the role of water inhe evolution of the Martian crust.

Direct measurements of water in Martian meteorites indicatehat the water contents of Martian magmas (e.g., McCubbint al., 2010) are much higher than the venerable Dreibus andänke (1987) estimate. McCubbin et al. (2010) estimate from

ydrous amphibole in melt inclusions in Chassigny that the mantleource region for the Chassigny magma contained between 130nd 250 ppm H2O, depending on the amount of partial melting

equired to produce the observed magma. Leshin (2000) measuredhe concentration of H2O in apatite in a depleted shergottite (QUE4201) and McCubbin et al. (2012) measured the concentrationf H2O in apatite an enriched one (Shergotty). From the observed

converted to H2O equivalent (near the surface, H could be bound as OH or H2O). Thecorrelation coefficient is substantially less than the 0.5 significance level requiredfor other element pairs, but nevertheless informative (see text).

concentrations and an estimate of when apatite crystallized in thelava, McCubbin et al. (2012) estimate that for these two samples,their parent magmas contained between 730 and 2870 ppm H2O(ignoring loss from the lavas). If they were produced by 10% partialmelting, their mantle source regions contained between 73 and287 ppm H2O. Hallis et al. (2012) measured H2O in apatite crystalsin the unweathered fall Nakhla, finding H2O contents similar tothose in Shergotty (0.46–0.64 wt%). This suggests a similar watercontent for the mantle source of the nakhlites; using the approachtaken by McCubbin et al. (2010) I estimate a mantle concentrationof 150–220 ppm. In contrast to these estimates, Usui et al. (2012)analyzed H2O in melt inclusions in olivine in Y980459. They findno evidence for degassing of the inclusions, yet find an average ofonly 146 ppm in the melt inclusions, hence a pre-degassing magmacontent with the same value. They calculate that this indicates amantle source containing 15–47 ppm H2O, in the range estimatedby Dreibus and Wänke (1987). One possible interpretation ofthese results is that water is heterogeneously distributed in themantle. In addition, the water content of the mantle is likely tohave decreased with time: the extensive fluvial alteration of thesurface attests to significant water release early in Martian history.Thus, the young igneous rocks on which these mantle waterestimates are based may reflect a degassed mantle that containedsignificantly more water initially.

An independent estimate can be made from the mean concen-trations of H2O and Cl on the surface of Mars as determined byGRS (Boynton et al., 2008). Both species behave incompatibly dur-ing partial melting and both degas from lava flows. If, in spiteof their complex behavior during weathering and other aqueousprocessing, they maintain an approximately constant ratio andboth concentrate to the same extent in the upper crust, we can esti-mate their abundance from the H2O/Cl ratio and our estimate bulkCl content of 32 ppm. GRS data (which measures H in the upper fewtens of centimeters) gives an equatorial (between ∼52.5 degreesnorth and south latitude) H2O concentration (water equivalent ofmeasured H) of 3.9 ± 1.9 wt% (2-� of total variation; standard errorof 1500 points gives a 2-� of the mean of 0.05 wt%). The large por-tions of the crust at higher and lower latitudes have very high H2Oconcentrations (e.g., Boynton et al., 2008), implying a much highersurface concentration, and higher H2O/Cl. Taking a very conserva-tive water concentration of 5 wt% for the entire surface gives H2O/Cl

of 10 and implies a bulk Mars concentration of 330 ± 10 ppm. As arough check, a plot of H2O versus Cl, when forced through zero(Fig. 15), gives a H2O/Cl ratio of 8.1 ± 0.5, and a bulk Mars H2O of260 ± 16 (2-�); however, the R2 is only 0.2. (The correlation line

er Erde 73 (2013) 401– 420 415

fiH

asmm21fi3ftoaeiF

scviiaW(wBDttoNat�(eeUYwe

6

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Fig. 16. Abundances in bulk silicate Mars normalized to CI chondrites for lithophile(top), chalcophile (middle), and highly siderophile (bottom) elements. Variationsin chalcophile elements probably reflect the extent to which the elements concen-trate in sulfide phases or their behavior when conditions force them to behave aslithophile elements. Except for Re, highly siderophile elements are present in CIrelative abundances (Brandon et al., 2012) and have abundances about a hundred

G.J. Taylor / Chemie d

t is essentially the same as using the average equatorial global2O/Cl ratio.)

Except for the interesting results from Y980459, meteorite datand the blatant evidence for extensive water action on the surfaceupport a relatively wet bulk composition for Mars. Geophysicalodeling results seem to support the idea of a wet mantle. Ther-al evolution models of crustal evolution (Hauck and Phillips,

002; Guest and Smrekar, 2005) require water abundances of00–1000 ppm to produce mantle viscosities low enough to allowor convection and ductile flow of mantle materials. Combin-ng all the data, I estimate that the Martian mantle contained00 ± 150 ppm of H2O, but considering the morphological recordor substantial surface water, the primitive mantle might have con-ained more than this amount. However, even a minimum valuef ∼300 ppm is about a factor of ten higher than the Dreibusnd Wänke (1987) estimate. Thus, their suggested mechanism fornriching FeO in Mars, reaction between H2O and metallic iron dur-ng accretion, was not effective enough to account for the observedeO concentration in bulk Mars.

The deuterium/hydrogen ratio might be diagnostic of theources of water to Mars (e.g., outer solar system objects such asomets versus inner solar system objects, adsorbed nebula waterersus water in phyllosilicates in planetesimals). The atmospheres enriched by a factor of 5 in D/H (�D value of a few thousand),mplying a large loss of water through sputtering by the solar windnd thermal escape of H preferentially to D (Owen et al., 1988).eathering products in Martian meteorites also have elevated D/H

�D of a few thousand). To understand the initial D/H in Martianater we need samples that have not been affected by atmosphere.ecause plate tectonics did not recycle the crust and its modified–H fractionated water, igneous rocks may contain the informa-

ion about the isotopic composition of primary Martian water, ifhey have not been altered after emplacement in the crust or whilen Earth. Hallis et al. (2012) measured H isotopes in apatite inakhla, a well-preserved fall. Terrestrial alteration is not detectablend Martian weathering is identifiable. Hallis et al. (2012) showhat the Nakhla parent magma had water with a terrestrial-likeD (−78 to +188); the terrestrial mantle has �D of −140 to +60Boettcher and O’Neil, 1980; Michael, 1988; Ahrens, 1989; Deloulet al., 1991; Bell and Rossman, 1992; Thompson, 1992; Grahamt al., 1994; Jambon, 1994; Wagner et al., 1996; Xia et al., 2002).sui et al. (2012) measured D/H in melt inclusions in olivine in980459, finding a �D of +275. It appears that bulk Mars beganith a D/H similar to that of Earth, though it could be slightly

levated.

. Discussion

The composition of bulk silicate Mars derived above is shownn Tables 4 and 5. Fig. 16 shows the results normalized to CI chon-rites (Lodders, 2003, except for highly siderophile elements where

use data for Orgueil, Horan et al., 2003), and plotted in orderf decreasing 50% condensation temperature (hence increasingolatility). Cl is plotted with Br and I, rather than its calculatedondensation temperature. The pattern for lithophile elementss the familiar one with uniform refractory element abundancesassumed, but reasonably so, Figs. 2 and 3). The exceptions are Fend P, where significant fractions were fractionated into the coreuring primary differentiation. Mars has roughly uniform depletionf moderately volatile elements (0.6 × CI), and strong depletion ofighly volatile elements. The highly volatile elements are within

ncertainties uniformly depleted at about 0.06 CI abundances. Theighly volatile chalcophile elements (Fig. 16) are likewise roughlyniformly depleted, but with more scatter. They have normalizedbundances of 0.03 × CI. I discuss these abundances in more detail

times higher than expected if they equilibrated with a metallic phase during coreformation.

below with the goal of showing the utility of knowing Martian (andother planetary) compositions.

The striking feature of the revised composition is how similarit is to that derived by Wänke and Dreibus, in spite of anal-yses of numerous newly found Martian meteorites, orbital andlander geochemical data, and improved global geophysical data.The robustness of the model stems from their geochemical insightabout element behavior, allowing us to determine volatile ele-ments from the assumed abundances of refractory elements andto assess bulk FeO, and the certainty that we have meteorites fromMars. It seems remarkable that the conclusions reached by Wänkeand Dreibus based on only ∼10 SNC meteorites stand the test of

abundant new data. The only significant difference between myreassessment and the original Wänke–Dreibus model is in the con-centration of H2O. Other differences are a matter of degree but

4 er Erde 73 (2013) 401– 420

ddaa

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ogheatw

cltbadpotiprwtOotiaatm

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ptKetcidrtBSm0os

atdotat

Fig. 17. CI-normalized abundances for Earth and Mars compared. Refractory ele-ments between Cr and Sr are shown for context; the full list of concentrations for

16 G.J. Taylor / Chemie d

o not change the basic picture: moderately volatile elements areepleted by a small amount compared to CI chondrites (factor ofbout 0.6), while highly volatile elements are depleted by a factorpproaching 50.

.1. Mars is rich in FeO

There is no doubt that the Martian mantle is much richer inxidized iron than is the terrestrial mantle. Both geochemical andeophysical data confirm it. One might argue that melting of aydrous mantle might have enriched magmas in FeO (Nekvasilt al., 2007), but Taylor et al. (2006a) marshal experimental data torgue that any magma generated by partial melting of a wet Mar-ian mantle and subsequent fractionation of that hydrous magmaould lead to lower FeO than in the original mantle peridotite.

Orbital measurements of Mercury (e.g., Peplowski et al., 2011)onfirm the suggestion (Robinson and Taylor, 2001) that it has aow FeO content (2–4 wt%). Robinson and Taylor (2001) argue onhe basis of surface analyses of Venus (Surkov et al., 1987) that itsulk FeO is close to that of Earth, about 8 wt% (e.g., McDonoughnd Sun, 1995). It is tempting to suggest on the basis of these fewata points (only four, though they represent 100% of the terrestriallanets!) that there is a gradient in the oxidation state through-ut the inner solar system. Another way of looking at this is thathe planetesimals accreting to form the terrestrial planets variedn their oxidation state (E chondrites to L chondrites, for exam-le): Mercury received a bigger share of the reduced chondriticaw materials than did Earth than did Mars. This is consistentith the modeling designed to match oxygen isotopic composi-

ions (Sanloup et al., 1999; Lodders and Fegley, 1998; Burbine and’Brien, 2004). Alternatively, perhaps the FeO was made on Mars byxidation of metallic iron. Wänke and Dreibus (1994) hypothesizedhat the higher FeO was caused by metal oxidation by H2O dur-ng the accretion of Mars. As discussed by Bertka and Fei (1998b),n oxidation event like this would increase the sulfur content of

metallic core. Whatever the cause, the different FeO concentra-ions of the terrestrial planets must be taken into account when

odeling planetary accretion.

.2. Depletion of volatile elements

Volatile elements fall into two distinct groups on the abundancelots (Fig. 16). One involves the moderately volatile (50% condensa-ion temperatures between ∼800 and 1100 K) lithophile elements, Ga, Na, Rb, Cs, and F, and the moderately volatile chalcophilelements As, Cu, and Ag. The lithophiles are depleted by small fac-ors compared to CI chondrites (the mean depletion is 0.6 × CI). Inontrast, the moderately volatile chalcophile elements are signif-cantly more depleted, average 0.03 × CI. This suggests that theirepletion is driven by core formation (probably S-rich, see below)ather than volatility. Highly volatile elements (50% condensationemperatures <750 K) are strongly depleted. Lithophile element (Cl,r, I) abundances are 0.06 × CI and chalcophile elements (Bi, Zn, Sn,e, Cd, In, and Tl) abundances are 0.03 × CI. The chalcophile ele-ents have a large range in depletion factors, from 0.004 (Se) to

.09 (In). This range likely reflects the combination of formationf a core rich in S and once sulfide was depleted in the mantleubsequent lithophile partitioning of the elements.

All the volatile elements are combined in Fig. 17, along with few refractory elements for reference, and compared to terres-rial abundances. Moderately volatile elements are somewhat moreepleted in Earth than in Mars (0.4 versus 0.6 × CI), but on average

verlap for the highly volatile elements. The depletion patterns forhe two planets track each other, discrepant most strongly for Brnd Tl. In spite of the difference for moderately volatile elements,he two bodies are quite similar. The biggest difference between

Mars is shown in Table 4. The overall pattern suggests similar compositions for Earthand Mars, though Mars contains somewhat higher concentrations of moderatelyvolatile elements.

Mars and Earth is that Mars has more than double the FeO con-centration and a sulfur-rich core (see below). This emphasizes thesimilarity among the terrestrial planets. They may have differentmixtures of the ingredients that delivered the refractory elements,but the sources for the volatile elements may have been quite simi-lar. The methods that use oxygen isotopes to derive the compositionof Mars only fail because they add too much of the volatile compo-nent. Perhaps the chondrite groups proposed had similar refractoryelement abundances, but accretion occurred mostly before thevolatiles condensed.

6.3. Water: abundance and source

Mars appears to have somewhat less H2O than does Earth,300 ± 150 versus 500 ppm; the terrestrial value is a minimum(Mottl et al., 2007). In spite of appearing to have less waterthan bulk Earth, Mars has higher concentrations of moderatelyvolatile elements and about the same concentrations of highlyvolatile elements. This is, of course, convoluted because most highlyvolatile elements are chalcophiles. If we consider only Cl, Br, andI, which are lithophile, then Mars appears to be enriched com-pared to Earth, 0.06 versus 0.02 × CI (Fig. 17). Assuming that thehighly volatile elements were delivered to accreting Mars in water-bearing planetesimals resembling carbonaceous chondrites, thenthose objects contained less water than did those accreting to Earth.I estimate that the planetesimals contributing the 0.06 × CI contri-bution of halogens to Mars would have contained on average only0.5 wt% H2O to produce a bulk Martian water content of 300 ppm.The volatile-bearing planetesimals used in constructing the Earthwould have contained about 3 wt% H2O to produce a bulk watercontent of 500 ppm while adding highly volatile elements to bringthe terrestrial inventory to only 0.02 × CI abundances.

The somewhat damp but otherwise volatile-rich planetesimalsare likely to have been formed from typical inner solar systemmaterials. Otherwise, the D/H of Martian water would not be so sim-ilar to the terrestrial value (Section 5.9). Except for Jupiter-familycomets, which have D/H similar to the Earth (Hartough et al., 2011),astronomical measurements of D/H in comets suggest that outersolar system materials have elevated D/H. Thus, it appears that theinner solar system objects were fed from the same source of H2O,or at least from sources with the same D/H ratio.

This analysis of total water in Mars assumes that the planet

did not lose water during accretion or afterwards. It is possible,however, that water may have been lost preferentially by impactheating during accretion (e.g., Bond et al., 2010), might have reactedwith metallic iron during accretion, accounting for the high FeO

er Erde 73 (2013) 401– 420 417

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6

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6

hc2ica43cincctfbpfar

amdcCttMcc

6

b

Table 6Comparisons for two important chemical parameters among inner solar systemobjects.

K/Th Refs FeO (wt%) Refs

Mercury 5200 a 3 f,gVenus ∼3000 b 8 fEarth 2900 c 8 cMoon 360 d 13 hMars 5300 e 18 i

(a) Peplowski et al. (2011); (b) Surkov et al. (1987); (c) McDonough and Sun (1995);

G.J. Taylor / Chemie d

n Mars (Wänke and Dreibus, 1994), or lost during magma oceanrystallization (Elkins-Tanton et al., 2005).

.4. Highly siderophile elements: implications for accretion

The significance of the concentrations of highly siderophile ele-ents in Mars Earth, and the Moon were reviewed in detail byalker (2009). New analyses and an updated discussion appear

n Brandon et al. (2012). The highly siderophile elements are gen-rally in chondritic relative abundances (except for Re) and areresent at 0.04 × CI concentrations (Fig. 16). The terrestrial prim-

tive upper mantle contains about twice this abundance (Beckert al., 2006). The abundances in Mars are quite depleted (0.04 × CI),ut if equilibrated with metallic iron during core formation theyould be another factor of 40–50 times lower because their low-ressure metal-silicate partition coefficients are around 10,000. Asxplained in detail by Walker (2009) and Brandon et al. (2012),hree explanations have been advanced to explain the surprisinglyigh concentrations of highly siderophile elements: (1) inefficientore formation; (2) equilibrium at high pressures where the parti-ion coefficients are lower; and (3) late addition of materials richn siderophile elements. As Brandon et al. (2012) explained, allf these ideas have flaws, but noted that late addition may behe most straightforward explanation. One problem that deserves

ore attention is how a chondritic ratio can be preserved duringagma ocean crystallization and subsequent mantle melting.

.5. Halogen concentration of the crust

If the bulk Cl abundance is as low as it appears (Section 5.6), theigh Cl in the uppermost surface has interesting implications forrustal evolution. Half the K in Mars is in the mantle (Taylor et al.,006a). This should apply to Cl as well because like K it is highly

ncompatible during igneous processing, implying that the crustalontribution to the total inventory is 16 ppm. Since the crust aver-ges 57 km thick (Wieczorek and Zuber, 2004), making up about.6 wt% of bulk silicate Mars, its mean Cl content should be around50 ppm. GRS and MER data show that the average surface Cl con-entration is about 5000 ppm. This indicates that Cl is concentratedn the uppermost crust. Assuming that the upper few meters areot exceptionally enriched, mass balances indicate that the entirerustal inventory of Cl could be confined to the upper ∼4 km. Iforrect, this suggests efficient aqueous transport of Cl from deep inhe crust to the surface, or continuous aqueous transport to the sur-ace during construction of the crust. This analysis is complicatedy the unknown extent to which Cl could have been lost in a gashase even at depth in the crust; such loss could also fractionate Clrom K. Nevertheless, to first order, it seems likely that Cl is system-tically concentrated toward the surface, illustrating the importantole of water in the crustal geochemistry of Mars.

It is interesting to consider the ramifications if Cl is much morebundant, as argued by Taylor et al. (2010) from the chondriticean Cl/K ratio of the surface and calculated similar 50% con-

ensation temperatures (Lodders, 2003). Considering that Cl/Br ishondritic and Cl/I close to it, this implies that all the halogens haveI-normalized abundances of 0.6. If the condensations tempera-ures of Br and I are as low as Lodders’ (2003) calculations indicate,hen all the highly volatile elements would have been present in

ars at about the same CI-normalized abundance, implying thatore formation reduced the abundances of the highly volatile chal-ophile elements from 0.6 to the observed ∼0.03.

.6. The core

The most reliable estimate of core composition may be the oney Khan and Connolly (2008) based on geophysical data. They find

(d) Warren and Wasson (1979), Warren (1989); (e) Taylor et al. (2006a); (f) Robinsonand Taylor (2001); (g) Nittler et al. (2011); (h) Taylor et al. (2006a); (i) Wänke andDreibus (e.g., 1988) and this paper.

a core radius of 1680 km and a composition (derived from calcu-lated core density) of 75–78 wt% Fe + Ni and 22–25 wt% S. Using thenew estimated Martian bulk composition, the relative amounts ofFe and FeO (total Fe, as for all refractory elements, is about 1.9 timesCI), assuming that essentially all the Ni and S are in the core (ignoressmall fractions in the silicate portion), and assuming an original Sabundance similar to the moderately volatile elements (0.6 × CI),I estimate a core composition of 78.6 wt% Fe + Ni and 21.4 wt% S.Better estimates await a determination of the core size by seismicmeasurements to be done by the Interior Exploration using Seis-mic Investigations, Geodesy and Heat Transport (InSight) mission,scheduled to be launched in 2016.

6.7. Comparing planet compositions

A central reason for determining planetary compositions is tocompare them to deduce variations in compositions, conditionsin the solar nebula and chemical uniformity of it, accretion pro-cesses (including the extent of mixing), and differentiation styles(e.g., floatation crust or not). Planetary scientists often emphasizedifferences among the terrestrial planets, but the similarities arestriking (Table 6). The close similarities in the D/H ratios of Mars,Earth, carbonaceous chondrites, and Jupiter-family comets sug-gest a common source of water-bearing material in the inner solarsystem (Alexander et al., 2012). K/Th (Table 6) varies among theterrestrial planets: Mercury, 5200 ± 1800 (Peplowski et al., 2011);Venus, ∼3000 (Surkov et al., 1987); Earth, 2900 (Jagoutz et al.,1979; McDonough and Sun, 1995; Taylor and McLennan, 2009);Mars, 5300 ± 220 (Taylor et al., 2006a). However, these variationsseem less significant when compared to the K/Th ratios of thecarbonaceous chondrites (19,000 – McDonough and Sun, 1995).The terrestrial planets are depleted in volatile elements comparedto carbonaceous chondrites, but the depletions are not correlatedwith distance from the Sun. Oxygen isotopes are distinctive amongplanets and meteorite groups (Mittlefehldt et al., 2008). The differ-ence between Earth (�17O of 0‰ by definition), and Mars (�17Oof + 0.25‰), is much smaller than the total range observed amongchondrites �17O of −4.3 to +2.5‰). Warren (2011) emphasizes thesimilarity among terrestral planets and differentiated meteoritescompared to carbonaceous chondrites in the isotopic compositionsof O, Cr, and Ti.

Some chemical parameters do vary directly with heliocentricdistance. Bulk silicate FeO (Table 6) increases from ∼3 wt% inMercury (Robinson and Taylor, 2001; Nittler et al., 2011), to 8 wt%in Earth (McDonough and Sun, 1995) and Venus (Robinson andTaylor, 2001) to 18 wt% in Mars. FeO and the size of the metalliccores are inversely correlated. These trends suggest a range inoxidation conditions correlated with heliocentric distance, in con-

trast to the apparent weak correlation with heliocentric distancefor D/H, K/Th, and oxygen isotopes. As these differences can beproduced by varying oxidation conditions, they do not suggestthe terrestrial planets were formed from fundamentally different

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18 G.J. Taylor / Chemie d

aterials. On the contrary, the broad chemical similarities amonghese planets indicate substantial mixing throughout the innerolar system during planet formation, as suggested by dynamicalodels (O’Brien et al., 2006; Walsh et al., 2011).

cknowledgements

I thank Klaus Keil for inviting me to write this review. The workeported here was initially inspired by my participation on thears Odyssey gamma-ray spectrometer team, under a NASA Par-

icipating Scientist grant (JPL 1241588), and further inspired byork done on a Mars Data Analysis grant (NNX07AV40G). All of

his work benefitted from valuable conversations with Heinrichänke and Gerlind Dreibus on central issues about the compo-

ition of Mars. Helpful and thorough reviews by Scott McLennannd Hap McSween improved the manuscript and sharpened somerguments significantly. Recent research was supported in party the National Aeronautics and Space Administration throughhe NASA Astrobiology Institute under Cooperative Agreement No.NA09DA77A.

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