high spectral resolution uv to near-ir observations of...

Icarus 191 (2007) 581–602www.elsevier.com/locate/icarus

High spectral resolution UV to near-IR observations of Marsusing HST/STIS

J.F. Bell III ∗, T.M. Ansty

Cornell University, Department of Astronomy, Ithaca, NY 14853-6801, USA

Received 2 February 2007; revised 15 May 2007

Available online 24 July 2007

Abstract

We acquired high spectral and spatial resolution hyperspectral imaging spectrometer observations of Mars from near-UV to near-IR wavelengths(∼300 to 1020 nm) using the STIS instrument on the Hubble Space Telescope during the 1999, 2001, and 2003 oppositions. The data sets havebeen calibrated to radiance factor (I/F ) and map-projected for comparison to each other and to other Mars remote sensing measurements. Wesearched for and (where detected) mapped a variety of iron-bearing mineral signatures within the data. The strong and smooth increase in I/F

from the near-UV to the visible that gives Mars its distinctive reddish color indicates that poorly crystalline ferric oxides dominate the spectralproperties of the high albedo regions (as well as many intermediate and low albedo regions), a result consistent with previous remote sensingstudies of Mars at these wavelengths. In the near-IR, low albedo regions with a negative spectral slope and/or a distinctive ∼900 nm absorptionfeature are consistent with, but not unique indicators of, the presence of high-Ca pyroxene or possibly olivine. Mixed ferric–ferrous mineralscould also be responsible for the ∼900 nm feature, especially in higher albedo regions with a stronger visible spectral slope. We searched for thepresence of several known diagnostic absorption features from the hydrated ferric sulfate mineral jarosite, but did not find any unique evidencefor its occurrence at the spatial scale of our observations. We identified a UV contrast reversal in some dark region spectra: at wavelengths shorterthan about 340 nm these regions are actually brighter than classical bright regions. This contrast reversal may be indicative of extremely “clean”low albedo surfaces having very little ferric dust contamination. Ratios between the same regions observed during the planet-encircling dust stormof 2001 and during much clearer atmospheric conditions in 2003 provide a good direct estimate of the UV to visible spectral characteristics ofairborne dust aerosols. These HST observations can help support the calibration of current and future Mars orbital UV to near-IR spectrometers,and they also provide a dramatic demonstration that even at the highest spatial resolution possible to achieve from the Earth, spectral variationson Mars at these wavelengths are subtle at best.© 2007 Elsevier Inc. All rights reserved.

Keywords: Mars; Spectroscopy; Ultraviolet observations; Hubble Space Telescope observations

1. Introduction and background

The geochemistry, mineralogy, and geomorphology of themartian surface preserve a record of surface–atmosphere inter-actions through time, just as they do on the Earth. However,unlike the Earth, the early martian geologic record may nothave been obliterated by plate tectonics and/or an extensivehydrologic cycle. Thus, the nature of the early martian envi-ronment may still be preserved in the planet’s present regolith.The key observational challenges to testing this hypothesis are

* Corresponding author.E-mail address: [email protected] (J.F. Bell III).

0019-1035/$ – see front matter © 2007 Elsevier Inc. All rights reserved.doi:10.1016/j.icarus.2007.05.019

finding unique geographic provinces and using diagnostic spec-troscopic discrimination methods that can allow these materialsto be uniquely detected and their abundances and distributionsto be quantified.

Mars is an iron-rich planet, and fortuitously, iron is themost spectrally-active cation in visible to near-IR (VNIR) re-mote sensing observations of planetary surfaces. Thus, remotesensing of Mars at solar reflectance wavelengths has the po-tential to reveal significant information about the distribution,mineralogy, crystallinity, oxidation state, and physical proper-ties of iron-bearing materials on its surface (e.g., Burns, 1993;Bell, 1996). Further, assessing and characterizing the inventoryof primary ferrous (Fe2+) iron-bearing minerals like pyroxene

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582 J.F. Bell III, T.M. Ansty / Icarus 191 (2007) 581–602

and olivine and their potential weathering products—secondaryferric (Fe3+) iron-bearing minerals like hematite, goethite, orjarosite, provides a window into the history of alteration andperhaps even past environmental conditions on the Red Planet.

Laboratory studies have shown that well-crystalline iron ox-ides, oxyhydroxides, and oxyhydroxysulfates have distinctivespectral slopes and absorption bands in the visible to near-IR (e.g., Sherman et al., 1982; Sherman and Waite, 1985;Morris et al., 1985). Of particular relevance to Mars surfaceobservations are: (a) the position of the near-UV O2− → Fe3+charge transfer absorption edge that gives the iron oxides (andMars) their distinctive red to yellow colors; (b) the amount ofspectral structure of the long-wavelength wing of the near-UVcharge transfer edge from about 400 to 700 nm, caused by aseries of Fe3+ ligand field and Fe3+ ↔ Fe3+ electronic pairtransitions that are sensitive to the degree of crystallinity; and(c) the strengths and positions of two Fe3+ ligand field transi-tions in the red to near-IR: one (6A1 → 4T2

4G) near 650 nmthat exhibits minor variations with mineralogy, and the sec-ond (6A1 → 4T1

4G) near 860 to 900 nm that shows diagnosticchanges as a function of mineralogy (e.g., Sherman et al., 1982;Sherman and Waite, 1985).

The general absence of specific spectral absorption bands inprevious ground based Mars 400 to 700 nm spectra led to theinterpretation that most of the surface Fe3+ exists within amor-phous or nanophase ferric oxide minerals (e.g., Singer, 1982;Morris et al., 1989, 1997; Bell, 1992; Bell et al., 1990, 1993).Telescopic detection of an 860 nm band and a 650 nm inflec-tion provided early evidence for well-crystalline (submicron tomicron grain size) so-called “red” hematite (αFe2O3) on someregions of Mars at the ∼5% abundance level (e.g., Morris andLauer, 1990; Bell et al., 1990). These features were mapped andinterpreted in more detail from Phobos-2 ISM orbital VNIRobservations at much higher spatial resolution (Bibring et al.,1990; Murchie et al., 1993), though for only a small percentageof the surface because of the early demise of that mission. Morerecent infrared orbital observations from the MGS/TES instru-ment enabled higher spatial resolution detection and mappingof even coarser-grained (tens of microns and larger) so-called“gray” hematite in a number of localized regions on the sur-face (Christensen et al., 2000a), including Meridiani Planum,which was eventually selected as the landing site for the MarsExploration Rover Opportunity based partly on the MGS/TESgray hematite discovery. Subtle shifts in the position of the(6A1 → 4T1

4G) ferric transition towards longer wavelengths(890 to 920 nm) also provided tentative but ambiguous evidencefor the presence of ferric oxyhydroxide phases like goethite(αFeOOH) on Mars in ISM data (e.g., Murchie et al., 1993;Geissler et al., 1993). A number of other ferric phases havebeen inferred from previous studies (e.g., Burns, 1980, 1987;Singer, 1982; Burns and Fisher, 1990; Bell, 1992; Orenbergand Handy, 1992; Morris et al., 1990, 1993; Bishop et al.,1995; Bishop and Murad, 1996), perhaps the most interestingof which, in hindsight, is the hydrated ferric sulfate jarosite[(K,Na,H3O)Fe3(SO4)2(OH)6], which was never observed di-rectly in remote sensing observations but was speculated tooccur on Mars based on geochemical arguments (Burns, 1987).

The subsequent in situ discovery of jarosite within small expo-sures of ferric- and sulfur-rich sedimentary outcrops in Merid-iani Planum by the Opportunity rover team (e.g., Klingelhöferet al., 2004) vindicated the original speculation, and continuesto motivate the remote sensing and in situ characterization ofiron-bearing minerals.

Where exposed to view in low albedo regions, the (pre-sumably bedrock) primary ferrous mineralogy of Mars ap-peared, based on ground based telescopic and orbital ISMand Mars Global Surveyor Thermal Emission Spectrometer(MGS/TES) observations, to be dominated by high-Ca pyrox-ene (e.g., Singer et al., 1979; Mustard et al., 1993; Bandfield,2002). These dark rocks and soils were observed to exhibit5% to 15% deep Fe2+ absorption bands near 950 to 1000 nm,characteristic of relatively unweathered pyroxenes. Laboratorystudies have shown that accurate data on pyroxene chemistry(Fe, Ca, Mg abundances) can be obtained by analyzing subtleshifts in the position of the so-called “1 micron” and “2 micron”absorption features (e.g., Adams, 1974; Cloutis and Gaffey,1991; Burns, 1993; Clark et al., 1993; Morris et al., 2000).For example, many orthopyroxenes (e.g., hypersthene, bronzite,and enstatite) absorb strongly between 900 and 1000 nm. High-calcium clinopyroxenes (e.g., diopside) possess an absorptionfeature at or beyond 1050 nm, as do most olivines. In addition tothe 1- and 2-µm absorption features, narrow absorption featuresat 506 nm and 548 nm resulting from spin-forbidden Fe2+ tran-sitions are present in laboratory spectra of orthopyroxenes (e.g.,Burns, 1993). These present a challenge to detection, however,as they exhibit much shallower band depth (less than 5%) thanthe 1-µm feature (Morris et al., 2000).

Comparisons of telescopic and ISM spectra with spectra andpetrology of the SNC meteorites (e.g., Singer and McSween,1993; McSween, 1985; Mustard and Sunshine, 1995) indicatedthat the ferrous mineralogy of the dark regions appears to beconsistent with a mixture of high-Ca and low-Ca pyroxeneswith similarities to some terrestrial komatiitic basalts. Bell etal. (1997a) reported evidence of pyroxene based on band depthmeasurements of multispectral HST/WFPC2 data. These tele-scopic and ISM results have been dramatically confirmed athigher spatial resolution and superior spatial coverage by VNIRimaging spectroscopic observations conducted by the OMEGAinstrument on the Mars Express orbiter (e.g., Bibring et al.,2005, 2006).

Olivine has been remotely inferred telescopically (e.g.,Huguenin, 1987) and directly measured from MGS/TES (e.g.,Christensen et al., 2000b; Hoefen et al., 2003) and OMEGA(e.g., Bibring et al., 2005; Mustard et al., 2005) spacecraft or-bital infrared observations. Olivine may be widely present onMars in localized deposits and has been observed in large-scalesubsurface layers of tens to hundreds of kilometers in extent(Christensen et al., 2003; Mustard et al., 2005). The mineralhas been identified in situ by the Mars Exploration Rover Spiritin Gusev crater within basaltic rocks (e.g., Morris et al., 2004;McSween et al., 2004). Olivines possess spectra that often re-semble those of high-calcium clinopyroxenes in the visible-NIRregion (through approximately 2.2 µm), and the associated ab-sorption wings are generally less steep than those of pyroxene

STIS Mars observations 583

features (Clark et al., 1993). Visible to near-IR measurementsmay provide some sensitivity to the presence of olivine, butin regions where pyroxene or crystalline ferric oxides also in-fluence the long-wavelength end of our spectra it may not bepossible to uniquely detect the presence of olivine in these mea-surements.

At shorter wavelengths, reflectance spectra of geologicalmaterials of importance to planetary science, such as mafic sili-cates, feldspars, and iron oxides/hydroxides, frequently exhibitnarrow absorption features in the ultraviolet region (<450 nm)due to a variety of electronic processes from a number of dif-ferent cations (e.g., Wagner et al., 1987; Cloutis, 2002). Giventhe electronic processes associated with these features, and byanalogy with the structural and compositional effects which cancause measurable variations in absorption band wavelength po-sitions, widths, and intensities in mafic silicates (Cloutis, 2002),ultraviolet spectroscopy has the potential to provide otherwiseunobtainable compositional information. Resolvable absorp-tion features appear in the ultraviolet region of mineral spectra,in spite of the low overall reflectance, due to the intense metal–oxygen charge transfers which dominate this spectral region(Wagner et al., 1987; McCormack et al., 2006).

In this paper we attempt to exploit what is known or inferredabout the spectral and physical characteristics of iron-bearingminerals to interpret high spectral resolution (hyperspectral)UV to short-wave near-IR imaging spectroscopic observationsof Mars acquired from the Hubble Space Telescope (HST) in1999, 2001, and 2003. These near-global observations have rel-atively low spatial resolution by the standards of Mars orbitalobservations, but some of them span unique and diagnosticparts of the spectrum that are not being studied from current orprevious spacecraft at Mars, and some of them include coverageduring a rare martian meteorological event—a planet-encirclingdust storm—that make them valuable data sets for the charac-terization of dust aerosol spectral properties.

2. Observations and data reduction

2.1. Observations

The data presented and analyzed here were collected dur-ing the Mars oppositions of 1999, 2001, and 2003 using theSpace Telescope Imaging Spectrograph (STIS) aboard HST(e.g., Woodgate et al., 1998) (Table 1). High spectral resolu-tion and regional-scale spatial resolution images of Mars weregathered during four sets of observations at each opposition,at intervals of 90◦ of martian central meridian longitude. Us-ing the 52 arcsec × 0.2 arcsec slit on the STIS, the telescopeacquired the martian disc and slewed the slit across it, collect-ing hundreds of lines of spectra at each observation (Fig. 1).Each line records spatial and spectral detail along the axes of atwo-dimensional array (Fig. 2). Spatial resolution varied signif-icantly between each opposition and to a lesser degree betweendifferent observations during the same opposition. Angular res-olution of the data was determined in the vertical axis (alongthe slit) by the STIS detector spacing (50.7 milliarcsec; KimQuijano et al., 2003) and, in the cross-slit direction of the scan,

Fig. 1. Schematic view of the HST STIS slit scanning approach used for theseobservations. Background image is an HST/WFPC2 673 nm image of Mars.

by the slit width (0.2 arcsec). Unfortunately, it was not possibleto time the observations to obtain true Nyquist-sampled cross-slit resolution. Thus, resolution cell size at Mars was roughly4 times greater horizontally (in the direction of the scan) thanvertically (along the axis of the slit).

In 1999, the observations were made on April 27, May 1,May 6, and May 7, using the G750L grism to collect data in thevisible/NIR spectral regions (527–1026 nm) at 0.48 nm spac-ing (Fig. 3). Spatial resolution near the sub-Earth point wasapproximately 84 km in the scan direction and 21 km alongthe slit. Viewing conditions were generally excellent (low Marsatmospheric dust opacity), and abundant surface detail could bedetected.

The 2001 and 2003 observations were designed to collectdata in the UV/Vis wavelengths (289–570 nm). It was hopedthat these data, collected in a portion of the spectrum and ata spectral resolution unmatched by any existing or plannedMars-orbiting sensors, could corroborate other measurementscollected in the visible and infrared regions on martian surfacemineralogy, and potentially provide some new constraints onthe UV optical properties of airborne dust (Table 2). The STISG430L grism was used in conjunction with the same slit used


Table 1HST STIS observations of Mars: 1999, 2001, 2003

UTdatea

Timeb

UTWavelengths(nm)

Diam.(arcsec)

SE lat.(◦)

SE lon.(◦)

Phaseangle (◦)

Ls

(◦)Resol.c

(km/pixel)PROGIDd

HST cycle 8 STIS data990427 19:53 STIS: 523.6–1026.6 at �λ = 0.48 16.2 18.97 45.43 2.76 130.54 84.1 × 21.3 8152, Bell990427 21:16 STIS: 523.6–1026.6 at �λ = 0.48 16.2 18.98 65.68 2.81 130.57 84.1 × 21.3 8152, Bell990427 22:53 STIS: 523.6–1026.6 at �λ = 0.48 16.2 18.99 89.34 2.86 130.60 84.1 × 21.3 8152, Bell990501 15:46 STIS: 523.6–1026.6 at �λ = 0.48 16.2 19.55 310.20 5.96 132.39 83.9 × 21.3 8152, Bell990501 17:08 STIS: 523.6–1026.6 at �λ = 0.48 16.2 19.56 330.20 6.01 132.42 83.9 × 21.3 8152, Bell990501 18:45 STIS: 523.6–1026.6 at �λ = 0.48 16.2 19.57 353.86 6.07 132.45 83.9 × 21.3 8152, Bell990506 13:26 STIS: 523.6–1026.6 at �λ = 0.48 16.1 20.27 232.25 10.06 134.77 84.2 × 21.4 8152, Bell990506 14:48 STIS: 523.6–1026.6 at �λ = 0.48 16.1 20.28 252.25 10.11 134.80 84.2 × 21.4 8152, Bell990506 16:24 STIS: 523.6–1026.6 at �λ = 0.48 16.1 20.28 275.66 10.17 134.83 84.2 × 21.4 8152, Bell990507 7:14 STIS: 523.6–1026.6 at �λ = 0.48 16.1 20.37 132.74 10.68 135.14 84.3 × 21.4 8152, Bell990507 8:31 STIS: 523.6–1026.6 at �λ = 0.48 16.1 20.38 151.52 10.72 135.16 84.3 × 21.4 8152, Bell990507 10:08 STIS: 523.6–1026.6 at �λ = 0.48 16.1 20.39 175.18 10.78 135.19 84.3 × 21.4 8152, Bell

HST cycle 10 STIS data010809 11:48 STIS: 270–590 at �λ = 0.27 15.9 5.87 197.52 38.11 210.96 85.4 × 21.7 9052, Bell010809 13:08 STIS: 270–590 at �λ = 0.27 15.9 5.87 217.01 38.13 211.00 85.4 × 21.7 9052, Bell010809 14:58 STIS: 270–590 at �λ = 0.27 15.9 5.86 243.79 38.16 211.05 85.4 × 21.7 9052, Bell010809 16:21 STIS: 270–590 at �λ = 0.27 15.9 5.85 264.00 38.18 211.08 85.4 × 21.7 9052, Bell010810 11:53 STIS: 270–590 at �λ = 0.27 15.8 5.74 189.39 38.45 211.58 86.2 × 21.8 9052, Bell010810 13:12 STIS: 270–590 at �λ = 0.27 15.8 5.73 208.63 38.47 211.61 86.2 × 21.8 9052, Bell010810 15:02 STIS: 270–590 at �λ = 0.27 15.8 5.72 235.41 38.50 211.66 86.2 × 21.8 9052, Bell010810 16:25 STIS: 270–590 at �λ = 0.27 15.8 5.71 255.62 38.52 211.69 86.2 × 21.8 9052, Bell010814 8:57 STIS: 270–590 at �λ = 0.27 15.3 5.17 109.05 39.68 213.95 88.9 × 22.5 9052, Bell010814 10:17 STIS: 270–590 at �λ = 0.27 15.3 5.16 128.53 39.70 213.98 88.9 × 22.5 9052, Bell010814 12:07 STIS: 270–590 at �λ = 0.27 15.3 5.15 155.31 39.72 214.03 88.9 × 22.5 9052, Bell010814 13:30 STIS: 270–590 at �λ = 0.27 15.3 5.14 175.52 39.74 214.07 88.9 × 22.5 9052, Bell010904 21:44 STIS: 270–590 at �λ = 0.27 12.9 0.83 96.84 44.25 227.30 105.3 × 26.7 9052, Bell010904 23:01 STIS: 270–590 at �λ = 0.27 12.9 0.82 115.58 44.26 227.33 105.3 × 26.7 9052, Bell010904 0:51 STIS: 270–590 at �λ = 0.27 12.9 0.80 142.35 44.27 227.38 105.3 × 26.7 9052, Bell010904 2:14 STIS: 270–590 at �λ = 0.27 12.9 0.79 162.55 44.28 227.41 105.3 × 26.7 9052, Bell

HST cycle 12 STIS data030821 4:54 STIS: 270–590 at �λ = 0.27 25.0 −18.97 103.08 8.32 245.13 54.4 × 13.8 9738, Bell030821 6:24 STIS: 270–590 at �λ = 0.27 25.0 −18.97 125.03 8.28 245.17 54.4 × 13.8 9738, Bell030821 8:01 STIS: 270–590 at �λ = 0.27 25.0 −18.96 148.68 8.23 245.21 54.4 × 13.8 9738, Bell030821 9:36 STIS: 270–590 at �λ = 0.27 25.0 −18.96 171.85 8.18 245.26 54.4 × 13.8 9738, Bell030822 6:31 STIS: 270–590 at �λ = 0.27 25.0 −18.93 117.89 7.58 245.81 54.3 × 13.8 9738, Bell030822 8:01 STIS: 270–590 at �λ = 0.27 25.0 −18.93 139.84 7.54 245.85 54.3 × 13.8 9738, Bell030822 9:37 STIS: 270–590 at �λ = 0.27 25.0 −18.93 163.25 7.49 245.89 54.3 × 13.8 9738, Bell030822 11:13 STIS: 270–590 at �λ = 0.27 25.0 −18.93 186.66 7.45 245.93 54.3 × 13.8 9738, Bell030827 21:01 STIS: 270–590 at �λ = 0.27 25.1 −18.80 285.89 4.91 249.37 54.1 × 13.7 9738, Bell030827 22:30 STIS: 270–590 at �λ = 0.27 25.1 −18.80 307.59 4.90 249.41 54.1 × 13.7 9738, Bell030828 0:06 STIS: 270–590 at �λ = 0.27 25.1 −18.79 331.00 4.89 249.45 54.1 × 13.7 9738, Bell030828 1:42 STIS: 270–590 at �λ = 0.27 25.1 −18.79 354.41 4.88 249.49 54.1 × 13.7 9738, Bell030828 22:38 STIS: 270–590 at �λ = 0.27 25.1 −18.78 300.72 4.83 250.04 54.1 × 13.7 9738, Bell030829 0:07 STIS: 270–590 at �λ = 0.27 25.1 −18.78 322.42 4.84 250.08 54.1 × 13.7 9738, Bell030829 1:43 STIS: 270–590 at �λ = 0.27 25.1 −18.78 345.83 4.84 250.13 54.1 × 13.7 9738, Bell030829 3:19 STIS: 270–590 at �λ = 0.27 25.1 −18.78 9.24 4.84 250.17 54.1 × 13.7 9738, Bell

a Read 030821 as August 21, 2003.b Time given as the start of the ∼25 to 50 min observing sequence.c Resolution is the best spatial resolution at the sub-Earth point for images using the STIS/CCD.d Space Telescope Science Institute Program Identification number and Principal Investigator, for HST data archive access.

in 1999 to capture the shorter wavelengths. Spectral resolutionwas approximately 0.274 nm in both 2001 and 2003. Spatialresolution near the sub-Earth point in 2001 was roughly 85–105 km in the cross-slit scan direction by 22–26 km along theslit axis, varying by date (Fig. 4). In 2003, the observationswere made during the closest martian opposition in recordedhistory, and so spatial resolution was an impressive (by Earth-

based standards) 54 km cross-slit and 14 km along the slit axis(Fig. 5).

Viewing conditions in 2003 were in low atmospheric dustopacity conditions, but in 2001 Mars was in the middle of a rareplanet-encircling dust storm (e.g., Smith et al., 2002). Most sur-face features were totally obscured by the airborne dust (Fig. 6).Comparison of the 2001 and 2003 data sets, acquired by the


Fig. 2. Example single-slit STIS G750L grating image from April 27, 1999 scan. The vertical axis is along-slit spatial information and the horizontal axis iswavelength. The slit at this time was oriented such that most of the scene covered high albedo terrain, transitioning to a dark region near the bottom of the slit. Darkvertical lines show the wavelengths of several prominent solar spectral absorption lines, and the high-frequency noise typical of long-wavelength fringing effectscan be seen on the right (longest wavelength) side of the image.

Fig. 3. Examples of 1999 HST STIS scans across Mars at 750 nm (top row) compared to HST WFPC2 enhanced color composites acquired close in time on thesame martian day (Bell, 2003).

same instrumentation and covering many of the same surfaceregions, thus provides an opportunity to compare a limited setof surface features two years apart under very different at-mospheric dust opacity conditions, enabling the derivation of ahigh spectral resolution data set on airborne dust in the UV/Viswindow.

In the 1999 collection, the STIS data cover nearly 92% ofthe martian surface. The 2001 data, collected during the globalstorm, covers approximately 75% of the surface. In 2003, po-sitioning of the telescope resulted in coverage of 64% of Mars;since few areas of the surface were imaged more than once, this

registered data set includes higher incidence or emission anglesat some points compared to the other collections.

2.2. Data reduction and map projection

Raw STIS images, which are 2-dimensional arrays of spatialinformation along one (along-slit) axis and spectral informa-tion in the cross-slit direction (Fig. 2), were calibrated usingthe standard CALSTIS data reduction and calibration pipeline(e.g., McGrath et al., 1997). Estimated uncertainties on theseradiances are 5% to 10% based on previous assessments of


Fig. 4. Examples of 2001 HST STIS scans across Mars at 500 nm.

Table 2UV–Vis absorption features in iron-bearing minerals relevant to Marsa

Bandcenter, nm

Origin Comments

225 Fe–O LMCTb Very strong, wing within STIS range250–270 Fe–O LMCT Very strong, wing within STIS range285–320 Fe3+ electronic transition Center varies among ferric minerals325 Fe3+ pair transition Strength varies among ferric minerals365–380 Fe3+ electronic transition Center varies among ferric minerals400 Fe3+ electronic transition Observed primarily in hematite435–444 Fe3+ electronic transition Strong jarosite band440–480 Fe3+ pair transition Center varies among ferric minerals480–530 Fe3+ electronic transition Center varies among ferric minerals

a Based on Sherman et al. (1982) and Sherman and Waite (1985).b Ligand to metal charge transfer.

HST calibration pipeline-derived radiances (e.g., Bell et al.,1997b). Derived radiances were then converted to radiance fac-tor or I/F (where I is the observed radiance from Mars andπF is the solar spectral irradiance at the top of the martian at-mosphere at the time of the observations; e.g., Wehrli, 1986;Hapke, 1993; Bell et al., 1997a). Uncertainties in derived I/F

values could arise from non-Lambertian scattering by the sur-face and/or atmosphere (Hapke, 1993), or from uncertaintiesin the estimated solar spectral irradiance and its convolutionover the STIS grating spectral profile. The latter error source isexpected to be quite small, but the former is difficult or impossi-ble to quantify. However, previous and ongoing experience withMars surface studies shows that at the scale of our observationsand at the low dust opacities of our 1999 STIS observations, theLambertian assumption implicit in our I/F derivation is a rea-sonable approximation (e.g., McEwen, 1991; Bell et al., 1997b;Soderblom et al., 2006). Thus, we would conservatively es-timate that the systematic error on the conversion of STIS-derived radiances to I/F values is approximately 10% for the1999 data set. Systematic uncertainties may be higher for the2001 and 2003 STIS data because of the higher atmosphericopacities; however, we are unable to quantify the detailed mag-nitude of these uncertainties with the data and other informationavailable.

Multiple calibrated 2-dimensional (along slit spatial infor-mation, wavelength) images from different positions along the

Fig. 5. Examples of 2003 HST STIS scans across Mars at 500 nm. The hori-zontal black areas are from occulting bars built into the STIS CCD field of viewfor stellar occultation studies.

slit scan across Mars were then merged to form a 3-dimensional(along slit spatial information, cross-slit spatial information,wavelength) hyperspectral image cube data set for each set ofobservations. In order to display the data in either a cylindri-cal or Mollweide projection map, the data were registered andwarped to a martian latitude/longitude grid using proceduressimilar to those described in Bell et al. (1997a) for HST WFPC2images of Mars. There were a few important differences in themap projection routines for these STIS data, however. First,because Mars rotated a significant fraction (∼10–20%) of the


Fig. 6. HST WFPC2 natural color composite images of Mars acquired within a few days to weeks of several of our 2001 STIS slit scans. The effect of the 2001planet-encircling dust storm on discrimination of surface albedo features is apparent.

slit width in the time that it took to acquire and read out eachslit position’s CCD image, each “spaghetti strip” of STIS datahad to be separately map projected. These separately-mappedslit images were then merged into a single map later, but oc-casional black strips of “missing” data can still be seen in thefull-resolution maps because of inconsistencies in the amountof time that it took to read out each slit strip image (as wellas delays for occasional STIS or HST onboard data housekeep-ing tasks). And second, the scanned images were not exactlyaligned with the martian polar axis, nor were they necessarilyaligned with each other across the four stepwise scans requiredfor coverage of the full disk. This necessitated rescaling the dataspatially to a square resolution cell prior to geographic registra-tion. The resulting Mollweide or cylindrical maps used for thisanalysis were resampled to 1◦/pixel and consist of 360 cells oflongitude and 180 cells of latitude. Thus, the best spatial resolu-tion at the equator was approximately 60 km in the final maps.

2.3. Spectral smoothing

These very high spectral resolution data sets, generallyhigher than needed in searching for diagnostic solid state fer-ric/ferrous mineral absorption features on Mars, also containboth random and systematic noise. In particular, dividing thesolar spectrum (Wehrli, 1986) from the observed Mars radi-ances produced some hysteresis artifacts at certain wavelengths.These artifacts arise from small differences in the spectral reso-lution and sampling of the solar and martian spectra, and areespecially apparent near strong and narrow solar absorptionlines. In the 2001 and 2003 UV–Vis data, noise was dominantat short wavelengths (below 350 nm in 2001, below 320 nm in

2003) and fairly constant at longer wavelengths. Random noiselevels were approximately ±0.005 I/F (these and other re-ported random noise values represent 1 standard deviation ofthe variance) in 2001 between 350 and 570 nm (Fig. 7, leftpanel, lower curve); the ratio of typical I/F across the brightregions of Mars to signal fluctuation due to noise was between6 and 30, as I/F increased from approximately 0.03 at 350nm to 0.15 at 570 nm. In the 2003 data set, noise between 320and 540 nm caused signal variations of ±0.004 I/F (Fig. 7,right panel, lower curve); at wavelengths longer than 540 nm,noise decreased in amplitude. The ratio of average I/F value tonoise variation in this data was between 8.8 and 30, improvingwith increasing wavelength. Noise in the 1999 Vis-NIR data in-creased from ±0.0025 I/F at 550 nm to ±0.005 at 1000 nm(Fig. 8, lower curve). The ratio of I/F to noise variance rangedbetween 22 and 60, depending on wavelength and surface re-flectance.

Random noise and/or solar spectrum calibration artifactsadded a high-frequency signal across the data sets, which ul-timately made any apparent single- or few-band feature verydifficult to discriminate from noise, obviating much of the ad-vantage of the very high spectral resolution. In order to tradeoff spectral resolution for a reduction in noise, both spectralsmoothing and band aggregation were evaluated to determinetheir effectiveness while retaining as much information contentas possible. Smoothing was accomplished using a boxcar av-eraging algorithm, specifying smoothing in only the spectraldimension. Band aggregation simply averaged values over anumber of adjacent bands.

Both methods were effective at reducing random noise andsolar spectral line hysteresis artifacts that were particularly no-


Fig. 7. (Left) I/F spectrum from a bright pixel within Arabia from our 2001STIS data: (bottom) and an 8-band aggregated, despiked spectrum from thesame pixel (top). (Right) I/F spectrum from a bright pixel within Arabia fromour 2003 STIS data: (bottom) and an 8-band aggregated, despiked spectrumfrom the same pixel (top). For clarity, the upper spectrum has been offset by+0.2 I/F units in both plots. The absolute photometric uncertainty of the de-rived I/F values is estimated at 5% to 10%, based on previous assessments ofHST calibration pipeline-derived radiances for Mars (e.g., Bell et al., 1997a).

Fig. 8. Noise reduction methods utilized in our 1999 STIS data. An unalteredspectrum from a bright pixel in Arabia (bottom); the same pixel after 7-bandboxcar smoothing (center); and after 8-band aggregation (top). Eight-band ag-gregate data was used for most data processing in this research.

ticeable in the UV data and lower band numbers of the 1999data set (Figs. 7 and 8). In fact, both resulted in an improvementin clarity for visualizing features that extended over severalbands, particularly when examining band ratio or band depth

images (described below). Neither aggregation nor smoothingeliminated spectral artifacts entirely unless the limits encom-passed so many bands (on the order of 12 nm or greater) as toresult in significant loss of detail and information content. Bandaggregation had the benefit of reducing file size significantly;original files had 1026 spectral bands, which could be reducedto 256 or 128 (averaging 4 or 8 bands respectively) with es-sentially no loss of usable information but a noticeable increasein signal-to-noise ratio (SNR). Most of the analysis describedhere was accomplished with 8-band aggregated data, resultingin an effective spectral resolution of about 2 nm in the 2001 and2003 data sets, and 4 nm in the 1999 Vis-NIR data, still morethan adequate for solid state mineral spectroscopy studies.

2.4. Long-wavelength fringing correction

The STIS CCD, like most CCDs, has a known issue withfringing at long wavelengths (greater than 750 nm; Kim Qui-jano et al., 2003). The silicon in the detector becomes progres-sively more transparent to light as wavelength increases; by 980nm, the amplitude of the fringing effects can reach 25% of sig-nal when using the G750L grating (Kimble et al., 1998). Use ofa contemporaneous flat-field calibration image can reduce themagnitude of the fringing to as little as 2% of spectral ampli-tude (Walsh et al., 1997), and this strategy was adopted for theseSTIS observations. Still, in the 1999 observations, effects cor-responding to residual fringing oscillations can be seen in theform of high- and low-frequency spectral variations at wave-lengths beginning at 870 nm and present in all data across themartian disk. These residual calibration artifacts take the formof regular waves of increasing amplitude, with minima at ap-proximately 900 and 990 nm and a maximum at 950 nm, as wellas rapidly rising apparent reflectance at the long-wavelengthend of the data beyond 1 µm (Fig. 9). The magnitude of thisresidual fringing error was approximately 3% of the signalamplitude, measured against a low order polynomial fit curveto the data. Previous spectral data of Mars acquired by otherhigh-spectral-resolution sensors at these wavelengths [for ex-ample, by the Imaging Spectrometer for Mars (ISM) instrumentaboard the Soviet Phobos-2 spacecraft, as well as ground basedtelescopic imaging spectrographs] do not show these kinds ofshapes as features of the martian surface in either high or lowalbedo areas (e.g., McCord et al., 1982; Bibring et al., 1990;Bell, 1992; Mustard et al., 1993; Mustard and Sunshine, 1995;Merényi et al., 1996). Their presence everywhere in the 1999STIS data, coupled with the lack of any known or plausible can-didate for a mineralogical or atmospheric explanation for thesekinds of features, strongly argues that they are artifacts of thesensor and calibration process that should not be mistaken fordepartures from previous Mars data.

After convincing ourselves that the artifacts above 870 nmwere fringing artifacts that are uniform across the scene ina relative sense, we developed a “bootstrap” algorithm to re-move these artifacts in a uniform way from the data set byscaling the long-wavelength STIS measurements to values de-rived from a different, trustworthy, data set. For this bootstrapmethod we used the merged ISM-telescopic average bright re-


Fig. 9. (Top) The spectra of the bright area in the vicinity of Olympus/Amazonisacquired by ISM (upper curve) in 1989 and STIS (lower curve) in 1999. Thedashed line is the ISM spectrum translated to equal the reflectance in the STISdata at approximately 840 nm. (Center) The correction vector produced by di-viding ISM spectrum values by STIS values at wavelengths longer than 840 nm.(Bottom) The resulting “Bright Mars” composite spectrum used in spectral ra-tio curves for comparison with terrestrial analogs to martian surface materials.

gion spectrum from Mustard and Bell (1994) to generate asingle “correction vector” that could be applied to all of thespectra in the 1999 STIS data set. We used the Mustard and Bell(1994) Olympus/Amazonis spectrum (Spot 41) for our com-parison spectrum, because we observed roughly the same highalbedo, presumably dust-covered, surface area in the 1999 STISdata. A STIS “Bright Mars” spectrum was derived by averaging∼20 pixels centered on the Olympus/Amazonis region (Fig. 9).The average I/F value at 839 nm for that region in the STISdata was 0.287, and the corresponding ISM-telescopic spec-trum’s I/F at 839 nm was 0.298, a difference of only about 4%.To build the correction vector, the ISM-telescopic spectrum wasscaled to the STIS average spectrum at 837 nm then divided bythe corresponding STIS I/F values. Finally, all spectra in the1999 STIS data set were adjusted in the range 837–1027 nmby multiplying all of the spectra by this one constant correctionvector. I/F values at wavelengths less than 837 nm were notmodified. It is important to point out that this scaling methodonly changes the absolute shape of I/F data longward of 837nm, but it does not change any of the relative region-to-regionspectral variations inherent in the data set.

2.5. Analysis tools

As a result of the relatively low spatial resolution (comparedto orbital spacecraft observations, for example), each pixel isunlikely to contain spectrally pure but spatially unresolved sur-face material (with the possible exception of relatively largeexpanses of CO2 and H2O ices at the polar caps). The surface ofMars is made up of many different minerals of differing com-position, grain size, and crystalline structure, so surface spectrashould not necessarily be expected to precisely match labora-tory references of pure samples. Nonetheless, it has been knownfor decades that the overall average nature of the non-ice sur-face spectrum of Mars is similar to the spectra of many kindsof ferric–iron bearing minerals (see review above). Spectral re-flectances typical of these minerals rise smoothly through theultraviolet into the red regions (600–700 nm), though the slopevaries from region to region on Mars and particularly betweenthe sharply demarcated bright and dark portions of the surface.Further, most ferric oxides are characterized by a reflectancepeak between 700 and 800 nm (e.g., Morris et al., 1985, 2000).The challenge in interpreting these and similar observations isto identify specific surface constituents through discernable anddiagnostic spectral features, and to map areas where these con-stituents are relatively more prevalent, despite the mixing ofseveral iron-bearing (or other) components in unknown propor-tions.

A key tool in searching for spectral fingerprints of spe-cific minerals making up only a fraction of the surface areain each pixel is measuring the band depth. In our band depthcalculations, a straight-line continuum is drawn between twopoints on a spectrum, and the difference between an observedvalue at some point between the endpoints of the continuumand the continuum value at that point is measured (e.g., Clarkand Roush, 1984; Bell and Crisp, 1993). This difference isexpressed as a fraction (here a percentage) of the continuumvalue, such that a band depth of zero is no band, positive banddepth indicates a concave (absorption) feature, and negativeband depth indicates a convex feature.

Band ratios can also provide a useful tool to search for min-erals based on diagnostic spectral shapes, as well as a way toenhance differences in surface reflectance in order to visuallyidentify faint surface details. In this research, we used band ra-tioing to look for detail in the ultraviolet and short-wave visiblespectra collected in 2001 and 2003; we divided the I/F valuesin one chosen band by those in another. A related technique wasused to search for sub-pixel surface constituents by dividingspectra in all pixels by the spectrum of a representative brightmartian region presumed to be dominated by ferric aeolian dust(cf. Merényi et al., 1996). While the absolute values of spectralratio curves shown here have little meaning when the ratios arebetween portions of the martian surface and laboratory spectraof reference materials, relative differences in absorption bandposition or spectral slope can be used to infer mineralogic sim-ilarities or differences. Thus, we scaled these curves to unityat 529 nm to facilitate comparison between laboratory mineralsamples and generally less reflective martian surface measure-ments.


We explored the possibility of using more complex analysismethods such as principal components analysis (PCA) or linearspectral mixture modeling on our STIS data. However, noneof these methods revealed useful new information relative tothe much simpler approaches outlined above. For example, un-raveling the physical meaning behind the eigenvectors derivedby PCA methods was problematic for these hyperspectral datasets. Also, it was not possible to untangle the effects of grainsize from abundance in simple linear mixing models at thesewavelengths. Thus, since our focus is primarily on relative re-gional spectral differences, our simpler ratio, band depth, andspectral slope methods were adequate to describe the variabil-ity that we observed in the data sets, as described below.

3. Results

3.1. General spectral trends and albedo features

The high iron content of the martian surface results in asteady positive spectral slope across the near-UV through thevisible region to about 600 nm, followed by a reflectancemaximum between 700 and 800 nm. In the UV–Vis, the re-sult is an average spectrum that is almost everywhere concave(slope gradually increases) from approximately 330 to 530 nm(Fig. 7), and convex from 530 to 800 nm (Fig. 8). Spatialvariations in this slope reveal interesting features despite thebroadly-similar nature of almost all surface spectra.

Most ferric oxide/oxyhydroxide minerals possess reflectancemaxima between 700 and 800 nm (e.g., Morris et al., 1985,2000). A reflectance peak appears as a negative value of banddepth when that parameter is defined relative to continuum val-ues on either side of the peak. For example, Fig. 10 shows a mapof the band depth at 759 nm defined using continuum points at701 and 818 nm, generated from our 1999 STIS data set. Whilevirtually all of Mars shows negative band depth as defined thisway at this wavelength, the areas most negative are those withthe highest albedo. Areas exhibiting less negative 759 nm banddepth values (brighter areas in Fig. 10) are lower albedo regionsthat also show evidence (as described below) for increased near-IR band depth attributed to ferrous iron-bearing minerals likepyroxene.

The spectral slope between about 530 and 570 nm in gen-eral correlates with albedo: the more reflective areas in thisrange show proportionally higher spectral slope than darker re-gions. Exceptions were observed in Mars’ western hemisphere,where Olympus Mons and the three volcanoes in the TharsisMontes possess notably lower spectral slope than surroundingregions of roughly similar albedo. For example, the ratio ofI/F at 570 nm to I/F at 530 nm for Olympus Mons is ap-proximately 1.35, and the value near the summit of AscraeusMons (the northernmost of the three large volcanoes in TharsisMontes) was 1.30. These values are significantly lower than thatin nearby high albedo terrain (∼1.47). For comparison, valuesfrom Syrtis Major and Arabia (classical low and high albedo re-gions, respectively) were 1.15 and 1.43, respectively. Radiancefactor values from the summit of Olympus Mons are higher thansurrounding terrain; at 657 nm the I/F value is 0.27 and at

Fig. 10. (a) Band depth at 759 nm, measured against a continuum between 701and 818 nm. Bright areas are those showing the least negative band depth (low-est absolute value), and dark areas the most negative band depth. The range ofthe image is −3.4% to −0.5% band depth. This is a Mollweide equal-area mapprojection of the data, with 90◦ N at the top, 90◦ S at the bottom, 0◦ longituderunning vertically down the middle, and 180◦ longitude along the peripheryof the map. (b) Global albedo map in the same projection, from MGS/TES(Christensen et al., 2001). (c) Global topography map in the same projection,from MGS/MOLA (Smith et al., 1999). White is high, black is low.

833 nm it reaches 0.31, among the highest values on the planetat those wavelengths. Spectral curves for the volcanoes appearsimilar to but translated to slightly higher values than the curvesfor surrounding terrain, at least partially explaining the lower570 to 530 nm band ratios.

Possible explanations for the higher I/F values over the vol-canoes include more reflective ice clouds near the summits,increased opacity of airborne dust with lower elevation, con-tributing to lower I/F at lower altitudes, or a layer of surfacefrost across the mountaintops. As the 1999 observations oc-curred during martian summer, the volcanoes are equatorialfeatures, and no typical water or CO2 ice spectra were observed,frost is an unlikely explanation. Likewise, the relative clarity ofsurface detail compared to observations made during periodsof higher airborne dust activity (Fig. 6), and the MGS/TES-


derived record of low atmospheric dust opacity during this time(e.g., Smith, 2004) argue against dust obscuration of lower al-titude terrain. Clouds are frequently observed in the vicinity ofMars’ higher peaks, especially during the aphelion season (e.g.,Wolff et al., 1999) and thus thin water ice clouds are the likelycause for these higher reflectance measurements.

A green to blue color ratio image (566/489 nm) from our2003 STIS observations reveals the presence of a subtle lin-ear albedo feature in the high albedo Arabia region (near 15◦N, 355◦ W) that does not appear to have been seen in pre-vious STIS or MOC data sets to correspond to any obvioussurface topographic or geologic feature (Fig. 11). The featureruns from northwest to southeast and does not appear to be anartifact of the data collection scheme, calibration, or the mapprojection, as it is roughly perpendicular to the orientation ofthe STIS slit and does not possess the smooth edges typicalof projection artifacts. Band ratio values (566/489 nm) withinthe feature average 1.805 ± 0.020, while northeast and south-west of the feature the values were 1.859 ± 0.010 and 1.847 ±0.013, respectively. The feature is barely discernable in the 8-band aggregated (2 nm spectral resolution) data at 566 nm, andthe pattern disappears entirely as wavelength decreases to ∼536nm. We searched for but could not find evidence of the fea-ture in MOC/WA high resolution images from 1999 (Fig. 11).Nor was the feature seen in a search of MOC/WA images from2003 around the time of our STIS observations (B. Cantor,personal communication, 2007). The lack of evidence for thefeature in MOC images and the lack of a correlation to anyobvious, permanent surface geologic feature suggest that thislow albedo marking may either be the result of small changesin the surface dust distribution in Arabia (perhaps as a resultof the planet-encircling dust storm of 2001), or that it may bethe result of a viewing geometry effect along a subtle albedoboundary that is simply more apparent in the specific late 2003

Fig. 11. Band ratio image (566:489 nm) of Mars from the 2003 STIS data (top,Mollweide projection). The dashed rectangle encloses the area enlarged at bot-tom right, showing the linear albedo feature in the Arabia region (arrow, at15◦ N, 355◦ W). At bottom left, a Mars Global Surveyor (MGS) Mars OrbitingCamera Wide Angle red filter “geodesy campaign” mosaic of the same region,generated from images acquired in 1999 (MSSS/JPL/NASA).

Earth-based viewing geometry than in nadir-viewing orbitalimaging.

3.2. Vis-NIR reflectance maxima (M1) and minima (T1)

As described above, martian surface reflectance spectra atSTIS spatial resolution are broadly similar, and show a max-imum value between 700 and 800 nm. Morris et al. (2000)examined the position of this reflectance maximum (termed“M1”) between 750 and 800 nm in Mars data from the Im-ager for Mars Pathfinder (IMP) instrument, comparing thoseresults to the M1 value for terrestrial ferric minerals and SNCmeteorites. They found that M1 could be the result of a com-bination of two or more superposed ferric mineral spectra, orthe interaction between a ferric absorption edge and the high-energy wing of a ferrous absorption feature. Because the 1999STIS data, like the IMP data set, do not extend significantly intothe 1 µm band region, this ambiguity precludes mineral identi-fication based solely on the location of this peak reflectance.However, the STIS data, with much greater spectral resolu-tion than IMP, allow more precise measurement of M1, andthe near-global coverage of the data set enables the trends tobe assessed as a function of albedo or other parameters. Asshown in Fig. 12, the position of M1 is linearly proportionalto the radiance factor at M1. This relationship holds acrossall data, whether from high or low albedo regions. The pres-ence of an M1 feature most likely implies the presence of atleast some crystalline ferric or ferrous minerals on the sur-face (Morris et al., 2000). The observed cluster of M1 valuesfor low albedo regions (∼725 to 750 nm) is consistent withthe presence of ferric-bearing minerals like hematite, goethite,and schwertmannite, and the cluster of M1 values for highalbedo regions (∼770 to 800 nm) is consistent with the presenceof ferric-bearing minerals like schwertmannite, lepidocrocite,maghemite, and ferrihydrite (e.g., Fig. 4 of Morris et al., 2000).

Fig. 12. The wavelength of the peak I/F value (M1) plotted against the I/F

value at that peak. Data points represent averages from 30 regions of interest en-compassing 1245 spectra from the 1999 STIS data set. Overall goodness-of-fitis 0.925. Separate regression curves for low albedo and high albedo subsets ofthe data are also shown. While there is more scatter in the fit to these subsets,the same overall correlation is still observed.


However, even with the higher spectral resolution of STIS,it is unlikely that the M1 parameter alone can be used to con-strain mineralogy at this level of detail. For example, Morriset al. (1995, 1997, 2000) have discussed how mixtures of fer-ric or ferric–ferrous minerals can lead to changes in peak re-flectances and band depths in visible to near-IR martian andterrestrial analog sample spectra. In particular, ferric mineral-ogy can be better constrained using both the M1 reflectancemaximum and the so-called “T1” reflectance minimum, whichoccurs between 750 and 1000 nm in many martian and lab-oratory mineral spectra. We measured the wavelengths of theM1 and T1 maxima and minima in the 1999 data by fittingspectra at each pixel to a sixth degree polynomial curve (inorder to minimize the effects of noise in determining minimaand maxima). M1 values were sought between 529 and 900nm, while T1 values were sought between 800 and 1020 nm.The resulting distribution of M1 values is very strongly cor-related with albedo, with classical bright regions exhibitinghigher values of M1 than classical dark areas. Based on com-parisons with laboratory studies (e.g., Morris et al., 2000;Cloutis and Bell, 2003), the observed M1 distribution is consis-tent with the Vis-NIR spectral properties of the martian surfacein the high albedo regions being dominated by ferric minerals,and the low albedo regions showing more evidence of ferrousminerals.

The distribution of T1 reflectance minima values is not asstrongly correlated with albedo (Fig. 13), and also appears tobe sensitive to minor variations in observed I/F caused by cal-ibration variance and different observation dates, resulting ina higher level of artifacts in the distribution map. Plotting T1against M1 as a histogram (Fig. 14) reveals a rather compactand correlated distribution of values. High values of M1, whichtend to correspond to low values of T1, are found in high albedoregions of Mars; conversely, low M1/high T1 points are fromthe darker portions of the surface. The range of M1/T1 val-ues plotted in Fig. 14 is consistent with the range derived forcandidate martian ferric and ferrous minerals by Morris et al.(2000). However, despite the near-global coverage of the dataset and the relatively high (by Earth-based standards) spatialresolution of the observations, our analysis does not reveal spe-cific regions of M1/T1 space that could be used to discriminateamong the specific kinds of ferric and ferrous minerals studiedby Morris et al. (2000). Rather, the correlated and somewhatbimodal nature of the M1/T1 distribution simply provide addi-tional support for the ferric-dominated nature of bright regionVis-NIR reflectances and the likelihood that ferrous mineralsare important Vis-NIR spectral components of the darker, pre-sumably bedrock areas.

3.3. Searching for jarosite

As noted above, jarosite is a ferric–iron bearing hydrated sul-fate of particular interest on Mars, as its presence had beenpredicted based on some Mars chemical weathering models(e.g., Burns, 1987) and ultimately it was discovered to makeup a small fraction (<10%) of the surface in scattered ar-eas of light-toned sulfur-rich sedimentary outcrops studied by

Fig. 13. Distribution of the T1 reflectance minimum, with brighter values indi-cating longer-wavelengths. The scale extends from 850 (black) to 970 (white)nm. Black areas within the image either did not possess a minimum between800 and 1021 nm or the minimum could not be reliably measured. Note that T1determination was highly sensitive to minor calibration variance; the diagonalbanding in the northern hemisphere is not correlated with known or observedsurface features.

Fig. 14. Two-dimensional histogram showing the correlation of measured M1I/F maxima and T1 I/F minima in our 1999 HST/STIS Mars data set. Darkerpixels correspond to higher histogram bin values.

the Opportunity rover (e.g., Klingelhöfer et al., 2004). One ofthe motivators of our STIS UV observations was that jarositepossesses a very characteristic 430–440 nm absorption bandwhich could be a key indicator of its presence even in ahighly mixed pixel (Fig. 15; Table 2; Sherman et al., 1982;Sherman and Waite, 1985; Morris et al., 2000). We searchedfor but could not detect any evidence of this feature in the 2003STIS data set using band ratios, spectral slope measurements,and band depth mapping.

Although the jarosite feature near 435 nm is unique and di-agnostic among ferric oxides, noise in the 2003 data set in therange between 400 and 440 nm is a significant obstacle to ob-serving this feature. The average I/F of the martian surfacerises from about 0.04 to 0.06 over that range (e.g., Fig. 24).Reference spectra of pure jarosite samples show a peak re-flectance of 0.14, surrounded by shoulders with approximately0.10 reflectance, for a peak magnitude of 1.4 times the base re-flectance (Fig. 15). Averaging I/F spectra across all of Marswithin this data set (8-band aggregate data, spectral resolutionapproximately 2 nm), the magnitude of fixed I/F noise arti-facts between 420 and 440 nm ranges between 0.002 and 0.01.This systematic noise occurs at fixed spectral values, and is


Fig. 15. Reflectance spectra (Morris et al., 2000) of iron-bearing minerals with relevance to Mars over the wavelength range of our STIS observations. Left: NarrowFe2+ absorptions in martian meteorite spectra. Right: Absorptions and/or inflections due to Fe3+ in a variety of iron oxides, oxyhydroxides, and oxyhydroxysulfates.See also Table 2.

present in roughly equal magnitude in all pixels; in addition,individual pixels exhibit random noise fluctuations on the or-der of 0.001 I/F . These noise sources result in an averagesignal-to-noise ratio of 40 to 60 (based on random noise) or 4to 6 (considering the systematic noise artifacts in the correcteddata). Expressed another way, the random noise magnitude isapproximately 4% of a baseline reflectance value of 0.05. Asthe jarosite peak is 40% greater than the baseline reflectance,jarosite abundance at the 10% level over broad areas of the mar-tian surface (comparable to the resolution of our data) would benecessary to be discerned at the same level as random noisein the data; much higher abundance levels would need to bepresent to exceed systematic noise artifact levels within our dataset. These abundances are not indicated either by in situ rovermeasurements nor other remote observations of the planet. Fur-thermore, band depth measurements across the jarosite featurewould show some spatially coherent patterns if the jarosite frac-tion were at approximately the same level as random noise; nosuch spatial patterns were observed.

Jarosite also exhibits a reflectance peak at 705 nm that couldbe an indicator of its presence (e.g., Morris et al., 2000). Un-fortunately, several other ferric minerals (e.g., hematite, schw-ertmannite, and akaganeite) also possess reflectance peaks inthe vicinity of 705 nm, so a conclusive search for jarosite mustexclude these other phases. While virtually all areas of Mars im-aged in the 1999 data set show positive band depth at 705 nm,no coherent spatial feature was found corresponding to a maxi-mum reflectance value at this wavelength within narrow searchbounds (∼20 nm). Additionally, the M1/T1 characteristics thatwould indicate discernable quantities of jarosite did not occurat surface locations where spectral evidence of jarosite could befound.

Thus, specific evidence of jarosite is not present in our 1999or 2003 UV to near-IR STIS data sets, either based on the435 nm or the 705 nm reflectance features. Seeking evidencefor jarosite in our airborne dust-dominated 2001 STIS UV dataset was no more fruitful. Band depth, band ratio, and spectralslope measurements generated from that data set in the appro-priate range showed only map-projection artifacts and low SNRartifacts.

3.4. Searching for pyroxene

As discussed above, the strength of the near-IR ferrous ab-sorption feature (“1 micron band”) can be used to assess theferrous iron-bearing mineralogy of the surface materials. Forexample, in our 1999 Vis-NIR STIS data, band depth at 900 nm,measured against a continuum between 741 and 1021 nm, cor-relates very strongly with the classical low albedo areas of Mars(Fig. 16). Our STIS 900 nm band depth map exhibits many sim-ilarities to the MEx/OMEGA equatorial to mid-latitude maps ofpyroxene derived by Bibring et al. (2005, 2006).

We generated similar band depth maps at other wavelengthsbetween 900 and 1000 nm, searching for evidence of slight spa-tial variations in band center position that could correlate with,for example, slight changes in pyroxene chemistry (Adams,1974; Cloutis and Gaffey, 1991; Bibring et al., 2005). Whilewe could detect differences in the absolute band depth values,we could not identify any statistically-significant spatial differ-ences in band depth maps made at these different band cen-ters. Unlike OMEGA spectra (e.g., Bibring et al., 2005, 2006;Mustard et al., 2005), our STIS data may not have high enoughSNR, or extend to long enough wavelengths, to allow us to de-tect subtle spatial variations in the martian spectrum in the 900to 1000 nm region. Thus, we use 900 nm band depth (Fig. 16)


Fig. 16. A band depth image of Mars, measuring depth at 900 nm as a fraction ofthe continuum value between 740 and 1021 nm. Areas in black represent banddepths <5.5%, blue represents band depths of 5.5–7.0%, cyan values representband depths of 7–9%, and yellow values represent band depths >9%.

as a proxy for a generic “900 to 1000 nm” (likely pyroxene)absorption feature on Mars.

Within the low albedo areas, we find significant and struc-tured variations in 900 nm band depth (Fig. 17). Syrtis Ma-jor and low albedo regions around Valles Marineris show thelargest band depth (approximately 13%), and the most wide-spread region of high band depth values occurs at the north-ern edge of the southern hemisphere low albedo regions. Lowalbedo northern hemisphere regions (besides Syrtis Major)show structured high 900 nm band depth as well, but maximumband depths are significantly lower than in Syrtis and southernhemisphere dark regions: about 8.4%. A circumpolar ring oflow albedo material at high northern latitudes possesses a high900 nm band depth (as high as 11%) relative to surroundingterrains. This dark ring corresponds to the north polar sand sea(e.g., Thomas and Weitz, 1989; Lancaster and Greeley, 1990).Similarly-high near-IR band depth values were reported in thisregion in previous HST multispectral imaging studies (e.g., Bellet al., 1997a, 1997b). In contrast, high albedo regions show 900nm band depth on the order of only 1–2%. The north polar capshowed very slight negative band depth, presumably a mani-festation of ices, rather than basaltic minerals, dominating thatregion’s near-IR spectral characteristics. Representative exam-ples of these spectra are shown in Fig. 17.

Spectra extracted from pixels showing the highest 900 nmband depth (Fig. 18, lower panel) show some general similari-ties to laboratory spectra of pyroxenes. I/F values in the high900 nm band depth spectra gradually rise through the near-UVand visible regions to a peak value of 11–13% near 750 nm,show a broad I/F minimum centered near 900 nm, then beginto increase again towards 1000 nm.

In an effort to identify whether pyroxene is indeed presentin these pixels, high 900 nm band depth spectra were dividedby the composite “Bright Mars” spectrum (Fig. 9) created whenbootstrapping the STIS data to the calibrated ISM observations.The resulting ratioed dark/bright spectra resembled neither lab-oratory spectra of pure pyroxene nor that of any specific ferrous(or ferric) mineral (Fig. 18, upper panel). The presence of aminimum value near 950 nm in some of these ratio spectra isconsistent with pyroxene, but could also be consistent with an-other mineral or mixture of minerals with an absorption feature

Fig. 17. Five spectra from our 1999 STIS near-IR data set, extracted from ar-eas representative of variations in ferrous iron-bearing mineralogy, and relevantto a search for pyroxenes (see text). At bottom (a), an averaged spectrum fromareas with greater than 5.5% band depth at 900 nm, measured against a con-tinuum between 740 and 1021 nm, which do not show increasing reflectancebetween 950 and 1005 nm. Next (b), an average spectrum from 10 selected re-gions (locations given in Table 3) also exhibiting high band depth at 900 nm,but with the more characteristic inflection near 900 nm. Spectrum (c) is fromthe north polar sand sea (22 pixel area, location approximately 75◦ N, 45◦ W),closely resembling the high band depth areas that were observed mostly in thesouthern hemisphere (see Fig. 16). Curve (d) is the reference “Bright Mars”spectrum, measured at Olympus Mons. At top, (e) is a typical spectral curvefrom the north polar cap (26 pixels, at approximately 80◦ N, 300◦ W). Spec-tra b, c and d are each offset by 0.05 from the curve below them; e is offset by0.20 from d. Representative 1σ error bars are shown at every ≈50 nm intervals.These uncertainties were derived by normalizing all spectra from each area at904 nm and then calculating the standard deviation of the normalized set.

near that wavelength (e.g., Morris et al., 1995, 1997, 2000). Themixed-pixel property of the STIS data, combined with Mars’ever-present aeolian dust, assure that a ferric reflectance maxi-mum in the vicinity of 750–800 nm will influence spectra every-where at this spatial resolution.

Burns (1993) noted the potential of 506 and 548 nm ab-sorption features as secondary identification features of Fe2+ inorthopyroxene laboratory spectra (Fig. 15). Unfortunately, nei-


Table 3Locations of the ten 3×3 pixel samples from within regions of high band depthat 900 nm against a continuum between 740 and 1021 nm

Region Latitude Longitude

N Polar Sand Sea 75◦ N 30◦ WN of Syrtis Major 53◦ N 285◦ WSyrtis Major 16◦ N 296◦ WSyrtis Major 8◦ N 290◦ WSyrtis Major 2◦ N 295◦ WS of Syrtis Major 12◦ S 292◦ WTerra Sabea 8◦ S 338◦ WMargaritifer Terra 5◦ S 20◦ WOphir Planum 6◦ S 52◦ WE End of Valles Marineris 18◦ S 60◦ W

ther of these shallow features were detected in either the 2001dust storm or 2003 clear atmosphere STIS observations thatspanned these wavelengths. The features are narrow enoughthat care had to be taken not to search too broadly, and band ag-gregated data of no more than 2 nm effective spectral resolutionwere used. No spatially coherent areas of positive band depthwere observed at either potential absorption feature wavelength.Sensors with greater spatial resolution and/or higher signal-to-noise ratios may be necessary to detect these subtle indicatorsof Fe2+in orthopyroxenes.

To further investigate the possible mix of ferric and fer-rous minerals present, reference spectra of two types of ferrousmineral laboratory spectra were divided by our STIS average“Bright Mars” spectrum (Fig. 18). The low-calcium pyroxene900–1000 nm absorption feature is prominent in the referencepyroxene to “Bright Mars” ratios, as well as in our actual STISspectra ratioed to “Bright Mars.” Substitution of palagonite forthe “Bright Mars” spectrum in spectral ratios with pyroxenesproduces curves which also possess minima consistent with thecurves produced using the martian spectrum (Fig. 19). Palag-onite [an amorphous mineraloid formed by hydration and de-vitrification of basaltic glass and which is pigmented to a red-dish color by nanometer-sized ferric oxide particles (np-Ox)dispersed throughout the hydrated basaltic glass matrix; e.g.,Morris et al., 1993] has been proposed as a martian surface ana-log material (e.g., Singer, 1982; Morris et al., 1989, 1993, 2000;Bell et al., 1993). Spectral ratios of pyroxenes to the hydratednp-Ox bearing mineral ferrihydrite also show similarities to py-roxene ratioed to the Bright Mars surface spectrum. Other morecrystalline (little or no np-Ox) samples of ferric minerals exam-ined (goethite, hematite, schwertmannite, akaganeite) producedcurves which were not similar. These results are consistent withthe spectral dominance of np-Ox on the martian surface at thesewavelengths, as opposed to coarser-grained, more crystallineferric or ferrous phases.

The pyroxene absorption feature may be detected between900 and 1100 nm, depending on the specific composition ofthe pyroxene variety. The percentages of iron, calcium, mag-nesium, chromium, and titanium in the sample, for example,can shift the wavelength at which this absorption feature falls(e.g., Adams, 1974; Cloutis and Gaffey, 1991). Additionally,olivines frequently possess a reflectance minimum at wave-lengths just past the long-wavelength end of our STIS Vis-NIR

Fig. 18. Top panel: Spectral ratio curves, showing hypersthene PYX02.h, en-statite NMNH 128288 (Clark et al., 1993), and the average spectrum from thesampled regions of high band depth at 900 nm. All spectra were divided byour “Bright Mars” spectrum. Note that all display a minimum between 900 and950 nm, characteristic of other pyroxene laboratory samples as well (Clark etal., 1993). Spectral ratios are scaled to unity at 529 nm. Bottom panel: The av-erage spectrum of the ten high band depth regions and those of hypersthenePYX02.h and enstatite NMNH 128288. The hypersthene spectrum has beenscaled (divided by 1.8) for comparison; that of enstatite has been divided by 5.

Fig. 19. Spectral ratio curves for ferrous minerals divided by ferric minerals thathave been proposed as Mars surface analogs: hypersthene (PYX02.h) dividedby palagonite (vol05a) and by ferrihydrite (GDS75), enstatite (NMNH128288)divided by palagonite and ferrihydrite, and the composite spectrum of the STIShigh band depth (900 nm) regions divided by the “Bright Mars” spectrum.Laboratory sample data from Singer (1982) and Clark et al. (1993). Curvesnormalized to their ratio value at 529 nm for comparison.


data (Fig. 20). Spectral differentiation of these minerals de-pends upon characterizing both this 900 to 1100 nm feature andthe associated 2000 nm absorption feature; the latter, however,is well outside the spectral range of our data. Large portions ofthe surface of Mars possess the high 900 nm band depth thatwould be expected for not only low-calcium pyroxenes but alsohigh-calcium pyroxene or sizeable surface deposits of olivine.Evidence supportive of either a high-calcium pyroxene or a va-riety of olivine within this data would consist of spectra whoseI/F values were not observed to rise longward of 950 nm, con-sistent with an absorption feature lying outside the observeddata range (Fig. 20).

3.5. UV-specific mineral and dust absorption features

McCormack et al. (2006) describe spectral features of lab-oratory reference materials in the near-UV indicative of min-erals that may be found on the martian surface. For example,olivine shows an absorption feature at 370 nm, montmorilloniteabsorbs at 370 and 395 nm, plagioclase feldspar absorbs at380 nm, and clinopyroxene, evidence for which also appears inthe NIR, absorbs at 395 nm. However, these features are all de-scribed as minor or weak, and none were observed in the 2001or 2003 data when using band depth as an investigative tool.Unfortunately, a major diagnostic feature for olivines and py-roxene that McCormack et al. (2006) illustrate at approximately300 nm falls in a region unusable in this data due to very highnoise.

There were some portions of the surface visible throughthe dust in the 2001 data. For example, Olympus Mons andthe Tharsis Montes were visible, particularly in longer wave-lengths, as slightly darker areas amidst the relatively bright ae-olian dust (Fig. 4). Valles Marineris, part of the seasonal southpolar cap, and a few other low albedo features could also bediscerned in the dusty 2001 STIS data by harshly stretchingthe image contrast at the longest wavelengths. However, thelevel of surface feature detail visible in 2001 is significantlydegraded compared to the 2003 data, which were collected un-der much lower atmospheric opacity conditions. Nonetheless,the appearance of recognizable features in limited portions ofboth the 2001 and 2003 data sets allows comparison of spectra,and in particular the characterization of UV–Vis spectra of dustaerosols from spectral ratios of the same regions in dusty andclear conditions (Figs. 21 and 22).

For example, measurement of average spectra within 3 × 3pixel regions over Olympus Mons, within the eastern portionof Valles Marineris, and within a relatively bright region northof Valles Marineris was possible in both the 2001 and 2003data. These spectra (Fig. 21) illustrate the contribution of theairborne dust, which is (not surprisingly) greater at lower alti-tude (the regions in or near Valles Marineris) than over featuresat much higher altitude, in this case a volcanic shield. The mea-sured I/F values at Olympus Mons are significantly lower thanboth lowland locations (Fig. 21, left panel), while in obser-vations much less affected by airborne dust (right panel) thedifference in I/F between the three regions is smaller. Whilegross surface detail can be seen through the dust in places, the

Fig. 20. Top panel: Spectral ratios of olivines GDS70.d and KI3005, and diop-side HS15 (Clark et al., 1993), divided by the “Bright Mars” spectrum, alongwith that of the averaged spectrum of the Mars regions whose spectral ratiocurves do not inflect upward between 951 and 1005 nm (heavy solid line) alsodivided by “Bright Mars.” Ratios are normalized to their values at 529 nm forcomparison. Lower panel: Spectra of the two olivines and diopside, comparedwith the averaged spectrum of the uninflected high band depth regions. Thelaboratory spectra are scaled for comparison; diopside and olivine GDS70.7spectra are divided by 4, and the olivine KI3005 spectrum is divided by 2.5.While the averaged spectrum from these areas of the martian surface does notstrongly resemble any of the reference olivines or clinopyroxenes whose char-acteristic absorption feature lies longward of 1000 nm, the ratio curves aresimilar in that they show negative slope more or less continuously across thespectral range of the STIS data.

contribution of surface reflectance to the observed I/F is lessthan that of the dust. This can be seen by comparing the spec-tra of two regions within and north of Valles Marineris in 2001and 2003. Both regions possess spectra that are very similar in2001, while in 2003 the spectrum of the higher-albedo regionto the north of the canyons exhibits higher I/F values at shortwavelengths and lower values at longer wavelengths than thosewithin Valles Marineris itself.

Spectral ratios support the greater contribution of airbornedust to the lowland spectra (Fig. 22). The spectral ratio fromOlympus Mons (solid line) possesses lower slope than the twofrom lower altitude; the poorly-crystalline ferric iron dominatedshape of the ratio curves is similar to that of the airborne dust


Fig. 21. UV–Vis spectra of the same Mars surface locations measured by STIS during the 2001 planet-encircling dust storm (left plot) and under low dust opacityconditions in 2003 (right). Olympus Mons, the eastern section of Valles Marineris, and a brighter region north of Valles Marineris are shown in averaged spectra of3×3 pixel areas. Note the significant difference in I/F values in the 2001 data between the lower-elevation regions and Olympus Mons; the reflectance contributionof airborne dust is much greater for the Valles Marineris observations. See text for details.

Fig. 22. Spectral ratios of 2001 data to 2003 data for the regions described inFig. 21: Olympus Mons, Valles Marineris, and the higher-albedo region to thenorth of Valles Marineris. The shape of all three ratios is similar to the airbornedust-dominated spectrum typical of the 2001 data set (see Fig. 7, for example).The ratio showing the least slope is that of Olympus Mons, which in 2001 wasless obscured by airborne dust than other locations on the planet’s surface.

itself, visible across the preponderance of the 2001 data (e.g.,Fig. 23). The high apparent noise in the ratio image is an am-plification of the relatively low SNR of the original UV spectra,particularly at shorter wavelengths.

While no other significant terrain types could be comparedbetween the 2001 and 2003 UV–Vis data sets due to limitedspatial overlap, data from 2001 partly covered the seasonalsouth polar cap and in fact this portion of the planet appeared

Fig. 23. Spectra from the 2001 data, from 3×3 pixel segments of the south polarcap region and an area at approximately 60◦ S, 140◦ W, a low-albedo region.These areas appear to be among the least affected by airborne dust in the 2001data. For contrast, spectra from two regions that appear to be highly affected byairborne dust (no surface detail visible at long wavelengths) are included. Theairborne dust spectrum appears globally similar, and dependent more on dateof collection than geographic location. The sample labeled “Dust–East” is fromapproximately 38◦ S, 253◦ W, while “Dust–West” is centered at 56◦ S, 78◦ W.

to be among those least obscured by dust during the 2001event. Fig. 23 shows a spectrum dominated by ice and/or thickclouds/haze in the south polar region, along with a spectrumfrom a relatively nearby low albedo area (in the vicinity of60◦ S, 140◦ W). These spectra are much more distinct fromeach other than the two regions in and near Valles Marineris,where surface features are less clearly visible through thedust.


Fig. 24. Sample spectra from the 2003 STIS data showing regions of varyingalbedo. All samples are 3 × 3 pixel regions; “Arabia” is centered at approxi-mately 2◦ N, 328◦ W; “Meridiani” at 5◦ S, 355◦ W; and “Syrtis” at 13◦ N,297◦ W. The low visible albedo Meridiani region exhibits an I/F contrastreversal compared to the higher albedo Arabia region at the shortest UV wave-lengths in our HST/STIS data set. The low visible albedo region Syrtis Majoralso shows a contrast reversal in the UV, but it is much smaller than that seenfor Meridiani. The inset shows an enlargement of the contrast reversal region.

Although the UV data obtained in 2003 did not allow detec-tion of fine features due to noise characteristics, spectral shapesfrom terrain across a range of classical albedos can be examined(Fig. 24). Higher-albedo areas such as Arabia Terra exhibit I/F

values that exceed those of Syrtis Major and Meridiani Planumat wavelengths greater than approximately 400 nm; brightnesscontrasts between bright and dark regions increase significantlyinto the visible and NIR portions of the spectra, where observedI/F ratios as great as ∼3:1 are seen between bright and darkareas.

At shorter wavelengths, the measured I/F of Mars is essen-tially flat between 300 and 380 nm, and begins to increase by400 nm. However, I/F values exhibit a shallow minimum of0.03 to 0.05 near 330 nm for a variety of terrain types in our2003 (clear) STIS data set. We observed a weak correlation be-tween high emission angles and I/F at 330 nm, suggesting thatatmospheric radiance (related to the presence of dust and iceaerosols) may play a role in determining the observed I/F forlimb and polar regions at these short wavelengths. Alternately,non-Lambertian surface scattering could also play a role, espe-cially at these more extreme viewing geometries. An interestingobservation was that the I/F differences between some non-limb, non-polar regions in the UV are not always correlatedwith their classical visible-wavelength albedos. For example,some classically bright regions appear to possess a lower I/F

(contrast reversal) at 300–380 nm than some dark regions. Thisobservation is discussed in more detail below.

4. Discussion and implications

4.1. Pyroxene on Mars

Cloutis and Gaffey (1991) used the positions of reflectanceminima at approximately 1 and 2 µm to constrain iron, magne-

sium and calcium fractions of laboratory pyroxene samples. Asthe 1999 NIR spectral data do not extend to the 2 µm feature,pyroxene cannot be uniquely identified on Mars based on theirmethods. In addition, the occurrence of substitution metals (Cr,Ti) in the pyroxene can shift reflectance minima in such a wayas to complicate identification based solely on spectral minima,even when spectra are unmixed. In the case of the STIS data, thepresence of some degree of ferric iron and the presence of unde-termined absorption features further casts into doubt the abilityto extract precise information on the type of pyroxene observed.Nonetheless, the general position of the feature between 940and 960 nm and its preferential occurrence in low-albedo re-gions does support the presence of a low-Ca clinopyroxene or amixture of clinopyroxenes and orthopyroxenes in some placeson Mars.

The regions that this method identified as likely possess-ing detectable deposits of pyroxene on the surface are largelyin agreement with other researchers’ findings. McCord et al.(1982) positively identified augite clinopyroxene as a con-stituent of Mars’ dark regions, using spectrographic NIR datacollected by a terrestrial telescope. Merényi et al. (1996) usedmultispectral visible-NIR data from ground based observationsand identified pyroxene distributions in Sinus Sabaeus and Si-nus Meridiani, immediately south of Arabia, as well as in Mar-garitifer Sinus. Gendrin et al. (2006) found pyroxene bandsto be strongest in Syrtis Major and in Valles Marineris, withslightly less strong indications in the southern low albedo re-gions. These findings were based on OMEGA data in the rangefrom 1.0 to 2.6 µm. Christensen et al. (2000b) used TES databetween 5.8 and 50 µm to map pyroxene from Mars orbit andagain found high concentrations in Syrtis Major and southernlow-albedo regions, particularly Tyrrhena Terra, dark terrainsouth of Elysium Planitia, Valles Marineris, and the fracturedand cratered terrain to its east. The combination of band depthmeasurements, band ratios producing curves very similar tothose of reference pyroxene spectra divided by an average spec-trum from high-albedo areas, and strong agreement with otherrecent findings using information in other spectral regions lendconfidence to identifying pyroxene as a surface constituent.

4.2. Olivine on Mars

Large surface deposits of olivine have been seen in a fewlocations, but not where our analysis would indicate. For exam-ple, Hoefen et al. (2003) detected an outcrop measuring 30,000km2 in the Nili Fossae region at the NE border of Syrtis Major,and Christensen et al. (2003) found olivine outcrops in GangesChasma. These have been based on longer-wave thermal in-frared observations and do not overlap the regions we identifiedin the 1999 STIS data. The combination of high band depthvalues at 900 nm and flat spectral reflectance values between951 and 1005 nm (factors most likely to be consistent withthe presence of olivine in our data set) were observed primar-ily in southern low albedo regions, including clusters of pixelsin the vicinity of 15◦ S, 280◦ W and 37◦ S, 232◦ W, and amuch more diffuse area southeast of Valles Marineris. High-calcium clinopyroxenes (such as diopside), like many olivines,


possess absorption features between 1000 and 1100 nm andthus decreasing reflectance between 900 and 1000 nm, and areplausible candidates for occurring mixed with other pyroxenesin the low-albedo regions on Mars. Unless higher spatial reso-lution data or longer-wavelength infrared observations providemore definitive confirmation of the presence of olivine in thesame areas where our 1999 HST/STIS analysis indicates thatit may occur, the conclusion that olivine is in fact present indiscernable quantities over these broad regions can only be con-sidered tentative.

4.3. UV reflectance contrast reversal

Fig. 24 shows that the I/F contrasts between some brightand dark regions on Mars reverse in the UV, below about340 nm. The effect appears to vary with region; the UV con-trast reversal observed in Meridiani is much stronger than thatobserved in Syrtis Major, for example, when compared to theUV I/F values of the bright (at visible wavelengths) regionArabia. A similar UV contrast reversal within the high northernlatitude low albedo polar sand sea was also recently noted byClancy et al. (2006) from initial Mars Color Imager (MARCI)UV observations on the Mars Reconnaissance Orbiter space-craft.

The strong UV contrast reversals seen in some low albedo re-gions may indicate extreme mineralogic differences in these re-gions compared to other areas on Mars. Specifically, the strongFe–O absorption feature that is partly responsible for givingMars its reddish color at visible wavelengths has a band min-imum at UV wavelengths that shifts slightly with variationsin Fe mineralogy and abundance (e.g., Wagner et al., 1987;McCormack et al., 2006). Thus, subtle changes in the Fe miner-alogy of some dark regions could cause the strong Fe–O absorp-tion to narrow, weaken, or shift to slightly longer wavelengths,thus increasing the UV reflectivity of these regions comparedto other parts of the planet. One possibility, consistent withstrong UV contrast reversals having been detected in some ofthe lowest (visible wavelength) albedo regions on the planet,is that the extremely “clean” (nearly dust-free) nature of theseregions removes or minimizes additional Fe–O UV absorptioncaused by ferric minerals in the dust and thus maximizes thedetectability of slightly different Fe–O absorption(s) from pri-marily ferrous minerals like pyroxene and/or olivine, or perhapsfrom some other iron-bearing mineral phase altogether that isnormally completely masked by ferric-rich dust in the UV.

Unfortunately, however, no systematic laboratory studieshave yet been performed attempting to correlate specific UVreflectance properties (band widths, depths, spectral shifts, etc.)with systematic variations in mineralogy for common ferric andferrous minerals of planetary interest, as has been done at visi-ble and near-IR reflectance wavelengths. Thus, it is not yet pos-sible to confirm whether or not a specific elemental or miner-alogic variation might be responsible for this contrast reversal.New laboratory studies being conducted on the UV reflectanceproperties of minerals and mineral mixtures (e.g., Cloutis etal., in preparation) will hopefully provide useful information to

help constrain the origin of the observed UV contrast reversalsseen in some martian dark regions.

5. Conclusions

• Measuring band depth at 900 nm, as a fraction of con-tinuum depth between 740 and 1021 nm, allows martianglobal pyroxene distribution to be mapped. More specifi-cally:– Low-calcium pyroxene, with an I/F minimum at wave-

lengths shorter than 1 µm, can be further characterizedby measuring the wavelength at which this minimumfalls, but our measurement is not diagnostic. Ferric min-erals affect the wavelength at which the minimum falls,and the requisite 2 µm feature falls outside the range ofour data.

– The presence of high-calcium pyroxene or surface de-posits of olivine is supported by spectra exhibiting high900 nm band depth but which appear to exhibit a re-flectance minimum at wavelengths well beyond 1 µm.

• STIS spectra in the Vis-NIR are consistent with the pres-ence of nanophase ferric oxide (np-Ox) as the dominantspectrally-active constituent of the martian surface at thesewavelengths. The spectral signature of np-Ox is seen every-where on Mars, regardless of albedo, at the spatial resolu-tion of these STIS measurements.

• Transport and removal of dust during the 2001 planet-encircling dust storm is proposed as a possible explana-tion of a significant “new” linear albedo feature observedthrough band ratio measurements in the Arabia high-albedoregion in 2003. This feature does not correspond to any ob-served long-term visible feature in the region; its temporalcharacteristics are unknown.

• Measurements of M1/T1 (near-IR reflectance maximum/minimum) were obtained and plotted for selected regionsfrom high- and low-albedo areas of Mars. While useful forclassifying portions of the surface in terms of similarity ordifference, the measurements do not provide unique resultsthat could be used for identifying specific surface minerals.Mixing of ferric minerals, or of ferric and ferrous minerals,moves the M1 feature away from wavelengths indicative ofpure ferric oxides. A bimodal distribution of M1/T1 val-ues that is correlated to albedo adds to the evidence for thespectral dominance of ferric chemistry in high albedo re-gions, and ferrous chemistry in low albedo regions.

• Comparison of UV/Vis spectra obtained in 2001 and 2003allows spectral characterization of aeolian dust in the 2001planet-encircling dust storm. Even during the 2001 storm,some surface detail could still be seen, particularly at highaltitudes (e.g., Olympus Mons) and at polar latitudes.

• Spectral contrast reversals seen in some dark regions in the2003 data at UV wavelengths below about 340 nm are con-sistent with recent MRO/MARCI observations, and may beindicative of the most dust-free low albedo martian surfaceareas.


Acknowledgments

This work is dedicated to the late Andy Lubenow, a plan-etary science colleague at STScI and an extraordinarily-skilledHubble Space Telescope observation planner. Andy devised andperfected the complex series of instrumental and spacecraftsequences that enabled us to obtain these unique and scien-tifically valuable data sets. We also thank Mike Wolff (SpaceScience Institute), Dick Morris (NASA/JSC), and Cornell stu-dents Jascha Sohl-Dickstein, Ben Cichy, and Brandi Wilcox fortheir assistance with the planning, processing, and map projec-tion of these data sets. We are grateful to G. Bellucci and ananonymous reviewer for their reviews of an earlier version ofthis paper. This work was funded by grants from the Space Tele-scope Science Institute (GO Programs 8152, 9052, and 9738)and Support for Program numbers GO8152, 9052, and 9738was provided by NASA through a grant from the Space Tele-scope Science Institute, which is operated by the Associationof Universities for Research in Astronomy, Incorporated, un-der NASA contract NAS5-26555. Crucial support has also beenprovided by NASA Planetary Geology and Geophysics pro-gram grant NNG04G163G.

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