High spectral resolution UV to near-IR observations of Mars using HST/STIS
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Icarus 191 (2007) 581–602
www.elsevier.com/locate/icarus
High spectral resolution UV to near-IR observations of Mars
using 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 have
been calibrated to radiance factor (I /F ) and map-projected for comparison to each other and to other Mars remote sensing measurements. We
searched 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 spectral
properties of the high albedo regions (as well as many intermediate and low albedo regions), a result consistent with previous remote sensing
studies of Mars at these wavelengths. In the near-IR, low albedo regions with a negative spectral slope and/or a distinctive ∼900 nm absorption
feature are consistent with, but not unique indicators of, the presence of high-Ca pyroxene or possibly olivine. Mixed ferric–ferrous minerals
could also be responsible for the ∼900 nm feature, especially in higher albedo regions with a stronger visible spectral slope. We searched for the
presence of several known diagnostic absorption features from the hydrated ferric sulfate mineral jarosite, but did not find any unique evidence
for its occurrence at the spatial scale of our observations. We identified a UV contrast reversal in some dark region spectra: at wavelengths shorter
than 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 storm
of 2001 and during much clearer atmospheric conditions in 2003 provide a good direct estimate of the UV to visible spectral characteristics of
airborne 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 variations
on 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 finding unique geographic provinces and using diagnostic spec-
troscopic discrimination methods that can allow these materials
The geochemistry, mineralogy, and geomorphology of the to be uniquely detected and their abundances and distributions
martian surface preserve a record of surface–atmosphere inter- to be quantified.
actions through time, just as they do on the Earth. However, Mars is an iron-rich planet, and fortuitously, iron is the
unlike the Earth, the early martian geologic record may not most spectrally-active cation in visible to near-IR (VNIR) re-
have been obliterated by plate tectonics and/or an extensive mote sensing observations of planetary surfaces. Thus, remote
hydrologic cycle. Thus, the nature of the early martian envi- sensing of Mars at solar reflectance wavelengths has the po-
ronment may still be preserved in the planet’s present regolith. tential to reveal significant information about the distribution,
The key observational challenges to testing this hypothesis are mineralogy, crystallinity, oxidation state, and physical proper-
ties of iron-bearing materials on its surface (e.g., Burns, 1993;
* Corresponding author. Bell, 1996). Further, assessing and characterizing the inventory
E-mail address: jfb8@cornell.edu (J.F. Bell III). of primary ferrous (Fe2+ ) iron-bearing minerals like pyroxene
0019-1035/$ – see front matter © 2007 Elsevier Inc. All rights reserved.
doi:10.1016/j.icarus.2007.05.019
582 J.F. Bell III, T.M. Ansty / Icarus 191 (2007) 581–602
and olivine and their potential weathering products—secondary The subsequent in situ discovery of jarosite within small expo-
ferric (Fe3+ ) iron-bearing minerals like hematite, goethite, or sures of ferric- and sulfur-rich sedimentary outcrops in Merid-
jarosite, provides a window into the history of alteration and iani Planum by the Opportunity rover team (e.g., Klingelhöfer
perhaps even past environmental conditions on the Red Planet. et al., 2004) vindicated the original speculation, and continues
Laboratory studies have shown that well-crystalline iron ox- to motivate the remote sensing and in situ characterization of
ides, oxyhydroxides, and oxyhydroxysulfates have distinctive iron-bearing minerals.
spectral slopes and absorption bands in the visible to near- Where exposed to view in low albedo regions, the (pre-
IR (e.g., Sherman et al., 1982; Sherman and Waite, 1985; sumably bedrock) primary ferrous mineralogy of Mars ap-
Morris et al., 1985). Of particular relevance to Mars surface peared, based on ground based telescopic and orbital ISM
observations are: (a) the position of the near-UV O2− → Fe3+ and Mars Global Surveyor Thermal Emission Spectrometer
charge transfer absorption edge that gives the iron oxides (and (MGS/TES) observations, to be dominated by high-Ca pyrox-
Mars) their distinctive red to yellow colors; (b) the amount of ene (e.g., Singer et al., 1979; Mustard et al., 1993; Bandfield,
spectral structure of the long-wavelength wing of the near-UV 2002). These dark rocks and soils were observed to exhibit
charge transfer edge from about 400 to 700 nm, caused by a 5% to 15% deep Fe2+ absorption bands near 950 to 1000 nm,
series of Fe3+ ligand field and Fe3+ ↔ Fe3+ electronic pair characteristic of relatively unweathered pyroxenes. Laboratory
transitions that are sensitive to the degree of crystallinity; and studies have shown that accurate data on pyroxene chemistry
(c) the strengths and positions of two Fe3+ ligand field transi- (Fe, Ca, Mg abundances) can be obtained by analyzing subtle
tions in the red to near-IR: one (6 A1 → 4 T2 4 G) near 650 nm shifts in the position of the so-called “1 micron” and “2 micron”
that exhibits minor variations with mineralogy, and the sec- absorption features (e.g., Adams, 1974; Cloutis and Gaffey,
ond (6 A1 → 4 T1 4 G) near 860 to 900 nm that shows diagnostic 1991; Burns, 1993; Clark et al., 1993; Morris et al., 2000).
changes as a function of mineralogy (e.g., Sherman et al., 1982; For example, many orthopyroxenes (e.g., hypersthene, bronzite,
Sherman and Waite, 1985). and enstatite) absorb strongly between 900 and 1000 nm. High-
The general absence of specific spectral absorption bands in calcium clinopyroxenes (e.g., diopside) possess an absorption
previous ground based Mars 400 to 700 nm spectra led to the feature at or beyond 1050 nm, as do most olivines. In addition to
interpretation that most of the surface Fe3+ exists within amor- the 1- and 2-µm absorption features, narrow absorption features
phous or nanophase ferric oxide minerals (e.g., Singer, 1982; at 506 nm and 548 nm resulting from spin-forbidden Fe2+ tran-
Morris et al., 1989, 1997; Bell, 1992; Bell et al., 1990, 1993). sitions are present in laboratory spectra of orthopyroxenes (e.g.,
Telescopic detection of an 860 nm band and a 650 nm inflec- Burns, 1993). These present a challenge to detection, however,
tion provided early evidence for well-crystalline (submicron to as they exhibit much shallower band depth (less than 5%) than
micron grain size) so-called “red” hematite (αFe2 O3 ) on some the 1-µm feature (Morris et al., 2000).
regions of Mars at the ∼5% abundance level (e.g., Morris and Comparisons of telescopic and ISM spectra with spectra and
Lauer, 1990; Bell et al., 1990). These features were mapped and petrology of the SNC meteorites (e.g., Singer and McSween,
interpreted in more detail from Phobos-2 ISM orbital VNIR 1993; McSween, 1985; Mustard and Sunshine, 1995) indicated
observations at much higher spatial resolution (Bibring et al., that the ferrous mineralogy of the dark regions appears to be
1990; Murchie et al., 1993), though for only a small percentage consistent with a mixture of high-Ca and low-Ca pyroxenes
of the surface because of the early demise of that mission. More with similarities to some terrestrial komatiitic basalts. Bell et
recent infrared orbital observations from the MGS/TES instru- al. (1997a) reported evidence of pyroxene based on band depth
ment enabled higher spatial resolution detection and mapping measurements of multispectral HST/WFPC2 data. These tele-
of even coarser-grained (tens of microns and larger) so-called scopic and ISM results have been dramatically confirmed at
“gray” hematite in a number of localized regions on the sur- higher spatial resolution and superior spatial coverage by VNIR
face (Christensen et al., 2000a), including Meridiani Planum, imaging spectroscopic observations conducted by the OMEGA
which was eventually selected as the landing site for the Mars instrument on the Mars Express orbiter (e.g., Bibring et al.,
Exploration Rover Opportunity based partly on the MGS/TES 2005, 2006).
gray hematite discovery. Subtle shifts in the position of the Olivine has been remotely inferred telescopically (e.g.,
(6 A1 → 4 T1 4 G) ferric transition towards longer wavelengths Huguenin, 1987) and directly measured from MGS/TES (e.g.,
(890 to 920 nm) also provided tentative but ambiguous evidence Christensen et al., 2000b; Hoefen et al., 2003) and OMEGA
for the presence of ferric oxyhydroxide phases like goethite (e.g., Bibring et al., 2005; Mustard et al., 2005) spacecraft or-
(αFeOOH) on Mars in ISM data (e.g., Murchie et al., 1993; bital infrared observations. Olivine may be widely present on
Geissler et al., 1993). A number of other ferric phases have Mars in localized deposits and has been observed in large-scale
been inferred from previous studies (e.g., Burns, 1980, 1987; subsurface layers of tens to hundreds of kilometers in extent
Singer, 1982; Burns and Fisher, 1990; Bell, 1992; Orenberg (Christensen et al., 2003; Mustard et al., 2005). The mineral
and Handy, 1992; Morris et al., 1990, 1993; Bishop et al., has been identified in situ by the Mars Exploration Rover Spirit
1995; Bishop and Murad, 1996), perhaps the most interesting in Gusev crater within basaltic rocks (e.g., Morris et al., 2004;
of which, in hindsight, is the hydrated ferric sulfate jarosite McSween et al., 2004). Olivines possess spectra that often re-
[(K,Na,H3 O)Fe3 (SO4 )2 (OH)6 ], which was never observed di- semble those of high-calcium clinopyroxenes in the visible-NIR
rectly in remote sensing observations but was speculated to region (through approximately 2.2 µm), and the associated ab-
occur on Mars based on geochemical arguments (Burns, 1987). sorption wings are generally less steep than those of pyroxene
STIS Mars observations 583
features (Clark et al., 1993). Visible to near-IR measurements
may provide some sensitivity to the presence of olivine, but
in regions where pyroxene or crystalline ferric oxides also in-
fluence the long-wavelength end of our spectra it may not be
possible to uniquely detect the presence of olivine in these mea-
surements.
At shorter wavelengths, reflectance spectra of geological
materials of importance to planetary science, such as mafic sili-
cates, feldspars, and iron oxides/hydroxides, frequently exhibit
narrow absorption features in the ultraviolet region (584 J.F. Bell III, T.M. Ansty / Icarus 191 (2007) 581–602 Table 1 HST STIS observations of Mars: 1999, 2001, 2003 UT Timeb Wavelengths Diam. SE lat. SE lon. Phase Ls Resol.c PROGIDd datea UT (nm) (arcsec) (◦ ) (◦ ) angle (◦ ) (◦ ) (km/pixel) HST cycle 8 STIS data 990427 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, Bell 990427 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, Bell 990427 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, Bell 990501 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, Bell 990501 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, Bell 990501 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, Bell 990506 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, Bell 990506 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, Bell 990506 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, Bell 990507 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, Bell 990507 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, Bell 990507 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 data 010809 11:48 STIS: 270–590 at λ = 0.27 15.9 5.87 197.52 38.11 210.96 85.4 × 21.7 9052, Bell 010809 13:08 STIS: 270–590 at λ = 0.27 15.9 5.87 217.01 38.13 211.00 85.4 × 21.7 9052, Bell 010809 14:58 STIS: 270–590 at λ = 0.27 15.9 5.86 243.79 38.16 211.05 85.4 × 21.7 9052, Bell 010809 16:21 STIS: 270–590 at λ = 0.27 15.9 5.85 264.00 38.18 211.08 85.4 × 21.7 9052, Bell 010810 11:53 STIS: 270–590 at λ = 0.27 15.8 5.74 189.39 38.45 211.58 86.2 × 21.8 9052, Bell 010810 13:12 STIS: 270–590 at λ = 0.27 15.8 5.73 208.63 38.47 211.61 86.2 × 21.8 9052, Bell 010810 15:02 STIS: 270–590 at λ = 0.27 15.8 5.72 235.41 38.50 211.66 86.2 × 21.8 9052, Bell 010810 16:25 STIS: 270–590 at λ = 0.27 15.8 5.71 255.62 38.52 211.69 86.2 × 21.8 9052, Bell 010814 8:57 STIS: 270–590 at λ = 0.27 15.3 5.17 109.05 39.68 213.95 88.9 × 22.5 9052, Bell 010814 10:17 STIS: 270–590 at λ = 0.27 15.3 5.16 128.53 39.70 213.98 88.9 × 22.5 9052, Bell 010814 12:07 STIS: 270–590 at λ = 0.27 15.3 5.15 155.31 39.72 214.03 88.9 × 22.5 9052, Bell 010814 13:30 STIS: 270–590 at λ = 0.27 15.3 5.14 175.52 39.74 214.07 88.9 × 22.5 9052, Bell 010904 21:44 STIS: 270–590 at λ = 0.27 12.9 0.83 96.84 44.25 227.30 105.3 × 26.7 9052, Bell 010904 23:01 STIS: 270–590 at λ = 0.27 12.9 0.82 115.58 44.26 227.33 105.3 × 26.7 9052, Bell 010904 0:51 STIS: 270–590 at λ = 0.27 12.9 0.80 142.35 44.27 227.38 105.3 × 26.7 9052, Bell 010904 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 data 030821 4:54 STIS: 270–590 at λ = 0.27 25.0 −18.97 103.08 8.32 245.13 54.4 × 13.8 9738, Bell 030821 6:24 STIS: 270–590 at λ = 0.27 25.0 −18.97 125.03 8.28 245.17 54.4 × 13.8 9738, Bell 030821 8:01 STIS: 270–590 at λ = 0.27 25.0 −18.96 148.68 8.23 245.21 54.4 × 13.8 9738, Bell 030821 9:36 STIS: 270–590 at λ = 0.27 25.0 −18.96 171.85 8.18 245.26 54.4 × 13.8 9738, Bell 030822 6:31 STIS: 270–590 at λ = 0.27 25.0 −18.93 117.89 7.58 245.81 54.3 × 13.8 9738, Bell 030822 8:01 STIS: 270–590 at λ = 0.27 25.0 −18.93 139.84 7.54 245.85 54.3 × 13.8 9738, Bell 030822 9:37 STIS: 270–590 at λ = 0.27 25.0 −18.93 163.25 7.49 245.89 54.3 × 13.8 9738, Bell 030822 11:13 STIS: 270–590 at λ = 0.27 25.0 −18.93 186.66 7.45 245.93 54.3 × 13.8 9738, Bell 030827 21:01 STIS: 270–590 at λ = 0.27 25.1 −18.80 285.89 4.91 249.37 54.1 × 13.7 9738, Bell 030827 22:30 STIS: 270–590 at λ = 0.27 25.1 −18.80 307.59 4.90 249.41 54.1 × 13.7 9738, Bell 030828 0:06 STIS: 270–590 at λ = 0.27 25.1 −18.79 331.00 4.89 249.45 54.1 × 13.7 9738, Bell 030828 1:42 STIS: 270–590 at λ = 0.27 25.1 −18.79 354.41 4.88 249.49 54.1 × 13.7 9738, Bell 030828 22:38 STIS: 270–590 at λ = 0.27 25.1 −18.78 300.72 4.83 250.04 54.1 × 13.7 9738, Bell 030829 0:07 STIS: 270–590 at λ = 0.27 25.1 −18.78 322.42 4.84 250.08 54.1 × 13.7 9738, Bell 030829 1:43 STIS: 270–590 at λ = 0.27 25.1 −18.78 345.83 4.84 250.13 54.1 × 13.7 9738, Bell 030829 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 resolution based standards) 54 km cross-slit and 14 km along the slit axis was approximately 0.274 nm in both 2001 and 2003. Spatial (Fig. 5). resolution near the sub-Earth point in 2001 was roughly 85– Viewing conditions in 2003 were in low atmospheric dust 105 km in the cross-slit scan direction by 22–26 km along the opacity conditions, but in 2001 Mars was in the middle of a rare slit axis, varying by date (Fig. 4). In 2003, the observations planet-encircling dust storm (e.g., Smith et al., 2002). Most sur- were made during the closest martian opposition in recorded face features were totally obscured by the airborne dust (Fig. 6). history, and so spatial resolution was an impressive (by Earth- Comparison of the 2001 and 2003 data sets, acquired by the
STIS Mars observations 585
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 is
wavelength. 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. Dark
vertical lines show the wavelengths of several prominent solar spectral absorption lines, and the high-frequency noise typical of long-wavelength fringing effects
can 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 the
same martian day (Bell, 2003).
same instrumentation and covering many of the same surface registered data set includes higher incidence or emission angles
regions, thus provides an opportunity to compare a limited set at some points compared to the other collections.
of surface features two years apart under very different at-
mospheric dust opacity conditions, enabling the derivation of a 2.2. Data reduction and map projection
high spectral resolution data set on airborne dust in the UV/Vis
window. Raw STIS images, which are 2-dimensional arrays of spatial
In the 1999 collection, the STIS data cover nearly 92% of information along one (along-slit) axis and spectral informa-
the martian surface. The 2001 data, collected during the global tion in the cross-slit direction (Fig. 2), were calibrated using
storm, covers approximately 75% of the surface. In 2003, po- the standard CALSTIS data reduction and calibration pipeline
sitioning of the telescope resulted in coverage of 64% of Mars; (e.g., McGrath et al., 1997). Estimated uncertainties on these
since few areas of the surface were imaged more than once, this radiances are 5% to 10% based on previous assessments of586 J.F. Bell III, T.M. Ansty / Icarus 191 (2007) 581–602
Fig. 4. Examples of 2001 HST STIS scans across Mars at 500 nm.
Table 2
UV–Vis absorption features in iron-bearing minerals relevant to Marsa
Band Origin Comments
center, nm
225 Fe–O LMCTb Very strong, wing within STIS range
250–270 Fe–O LMCT Very strong, wing within STIS range
285–320 Fe3+ electronic transition Center varies among ferric minerals
325 Fe3+ pair transition Strength varies among ferric minerals
365–380 Fe3+ electronic transition Center varies among ferric minerals
400 Fe3+ electronic transition Observed primarily in hematite
435–444 Fe3+ electronic transition Strong jarosite band
440–480 Fe3+ pair transition Center varies among ferric minerals
480–530 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 and
πF 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 uncertainties
in the estimated solar spectral irradiance and its convolution
over the STIS grating spectral profile. The latter error source is
expected to be quite small, but the former is difficult or impossi-
ble to quantify. However, previous and ongoing experience with
Mars surface studies shows that at the scale of our observations 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 view
and at the low dust opacities of our 1999 STIS observations, the
for stellar occultation studies.
Lambertian assumption implicit in our I /F derivation is a rea-
sonable approximation (e.g., McEwen, 1991; Bell et al., 1997b;
slit scan across Mars were then merged to form a 3-dimensional
Soderblom et al., 2006). Thus, we would conservatively es-
(along slit spatial information, cross-slit spatial information,
timate that the systematic error on the conversion of STIS-
derived radiances to I /F values is approximately 10% for the wavelength) hyperspectral image cube data set for each set of
1999 data set. Systematic uncertainties may be higher for the observations. In order to display the data in either a cylindri-
2001 and 2003 STIS data because of the higher atmospheric cal or Mollweide projection map, the data were registered and
opacities; however, we are unable to quantify the detailed mag- warped to a martian latitude/longitude grid using procedures
nitude of these uncertainties with the data and other information similar to those described in Bell et al. (1997a) for HST WFPC2
available. images of Mars. There were a few important differences in the
Multiple calibrated 2-dimensional (along slit spatial infor- map projection routines for these STIS data, however. First,
mation, wavelength) images from different positions along the because Mars rotated a significant fraction (∼10–20%) of theSTIS Mars observations 587
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 2001
planet-encircling dust storm on discrimination of surface albedo features is apparent.
slit width in the time that it took to acquire and read out each 2003) and fairly constant at longer wavelengths. Random noise
slit position’s CCD image, each “spaghetti strip” of STIS data levels were approximately ±0.005 I /F (these and other re-
had to be separately map projected. These separately-mapped ported random noise values represent 1 standard deviation of
slit images were then merged into a single map later, but oc- the variance) in 2001 between 350 and 570 nm (Fig. 7, left
casional black strips of “missing” data can still be seen in the panel, lower curve); the ratio of typical I /F across the bright
full-resolution maps because of inconsistencies in the amount regions of Mars to signal fluctuation due to noise was between
of time that it took to read out each slit strip image (as well 6 and 30, as I /F increased from approximately 0.03 at 350
as delays for occasional STIS or HST onboard data housekeep- nm to 0.15 at 570 nm. In the 2003 data set, noise between 320
ing tasks). And second, the scanned images were not exactly and 540 nm caused signal variations of ±0.004 I /F (Fig. 7,
aligned with the martian polar axis, nor were they necessarily right panel, lower curve); at wavelengths longer than 540 nm,
aligned with each other across the four stepwise scans required noise decreased in amplitude. The ratio of average I /F value to
for coverage of the full disk. This necessitated rescaling the data noise variation in this data was between 8.8 and 30, improving
spatially to a square resolution cell prior to geographic registra- with increasing wavelength. Noise in the 1999 Vis-NIR data in-
tion. The resulting Mollweide or cylindrical maps used for this creased from ±0.0025 I /F at 550 nm to ±0.005 at 1000 nm
analysis were resampled to 1◦ /pixel and consist of 360 cells of (Fig. 8, lower curve). The ratio of I /F to noise variance ranged
longitude and 180 cells of latitude. Thus, the best spatial resolu- between 22 and 60, depending on wavelength and surface re-
tion at the equator was approximately 60 km in the final maps. flectance.
Random noise and/or solar spectrum calibration artifacts
2.3. Spectral smoothing added a high-frequency signal across the data sets, which ul-
timately made any apparent single- or few-band feature very
These very high spectral resolution data sets, generally difficult to discriminate from noise, obviating much of the ad-
higher than needed in searching for diagnostic solid state fer- vantage of the very high spectral resolution. In order to trade
ric/ferrous mineral absorption features on Mars, also contain off spectral resolution for a reduction in noise, both spectral
both random and systematic noise. In particular, dividing the smoothing and band aggregation were evaluated to determine
solar spectrum (Wehrli, 1986) from the observed Mars radi- their effectiveness while retaining as much information content
ances produced some hysteresis artifacts at certain wavelengths. as possible. Smoothing was accomplished using a boxcar av-
These artifacts arise from small differences in the spectral reso- eraging algorithm, specifying smoothing in only the spectral
lution and sampling of the solar and martian spectra, and are dimension. Band aggregation simply averaged values over a
especially apparent near strong and narrow solar absorption number of adjacent bands.
lines. In the 2001 and 2003 UV–Vis data, noise was dominant Both methods were effective at reducing random noise and
at short wavelengths (below 350 nm in 2001, below 320 nm in solar spectral line hysteresis artifacts that were particularly no-588 J.F. Bell III, T.M. Ansty / Icarus 191 (2007) 581–602
images (described below). Neither aggregation nor smoothing
eliminated spectral artifacts entirely unless the limits encom-
passed so many bands (on the order of 12 nm or greater) as to
result in significant loss of detail and information content. Band
aggregation had the benefit of reducing file size significantly;
original files had 1026 spectral bands, which could be reduced
to 256 or 128 (averaging 4 or 8 bands respectively) with es-
sentially no loss of usable information but a noticeable increase
in signal-to-noise ratio (SNR). Most of the analysis described
here was accomplished with 8-band aggregated data, resulting
in an effective spectral resolution of about 2 nm in the 2001 and
2003 data sets, and 4 nm in the 1999 Vis-NIR data, still more
than adequate for solid state mineral spectroscopy studies.
2.4. Long-wavelength fringing correction
Fig. 7. (Left) I /F spectrum from a bright pixel within Arabia from our 2001 The STIS CCD, like most CCDs, has a known issue with
STIS data: (bottom) and an 8-band aggregated, despiked spectrum from the fringing at long wavelengths (greater than 750 nm; Kim Qui-
same pixel (top). (Right) I /F spectrum from a bright pixel within Arabia from
our 2003 STIS data: (bottom) and an 8-band aggregated, despiked spectrum
jano et al., 2003). The silicon in the detector becomes progres-
from the same pixel (top). For clarity, the upper spectrum has been offset by sively more transparent to light as wavelength increases; by 980
+0.2 I /F units in both plots. The absolute photometric uncertainty of the de- nm, the amplitude of the fringing effects can reach 25% of sig-
rived I /F values is estimated at 5% to 10%, based on previous assessments of nal when using the G750L grating (Kimble et al., 1998). Use of
HST calibration pipeline-derived radiances for Mars (e.g., Bell et al., 1997a). a contemporaneous flat-field calibration image can reduce the
magnitude of the fringing to as little as 2% of spectral ampli-
tude (Walsh et al., 1997), and this strategy was adopted for these
STIS observations. Still, in the 1999 observations, effects cor-
responding to residual fringing oscillations can be seen in the
form of high- and low-frequency spectral variations at wave-
lengths beginning at 870 nm and present in all data across the
martian disk. These residual calibration artifacts take the form
of regular waves of increasing amplitude, with minima at ap-
proximately 900 and 990 nm and a maximum at 950 nm, as well
as rapidly rising apparent reflectance at the long-wavelength
end of the data beyond 1 µm (Fig. 9). The magnitude of this
residual fringing error was approximately 3% of the signal
amplitude, measured against a low order polynomial fit curve
to the data. Previous spectral data of Mars acquired by other
high-spectral-resolution sensors at these wavelengths [for ex-
ample, by the Imaging Spectrometer for Mars (ISM) instrument
aboard the Soviet Phobos-2 spacecraft, as well as ground based
telescopic imaging spectrographs] do not show these kinds of
shapes as features of the martian surface in either high or low
albedo 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 1999
STIS data, coupled with the lack of any known or plausible can-
didate for a mineralogical or atmospheric explanation for these
kinds of features, strongly argues that they are artifacts of the
Fig. 8. Noise reduction methods utilized in our 1999 STIS data. An unaltered sensor and calibration process that should not be mistaken for
spectrum from a bright pixel in Arabia (bottom); the same pixel after 7-band departures from previous Mars data.
boxcar smoothing (center); and after 8-band aggregation (top). Eight-band ag- After convincing ourselves that the artifacts above 870 nm
gregate data was used for most data processing in this research.
were fringing artifacts that are uniform across the scene in
a relative sense, we developed a “bootstrap” algorithm to re-
ticeable in the UV data and lower band numbers of the 1999 move these artifacts in a uniform way from the data set by
data set (Figs. 7 and 8). In fact, both resulted in an improvement scaling the long-wavelength STIS measurements to values de-
in clarity for visualizing features that extended over several rived from a different, trustworthy, data set. For this bootstrap
bands, particularly when examining band ratio or band depth method we used the merged ISM-telescopic average bright re-STIS Mars observations 589
2.5. Analysis tools
As a result of the relatively low spatial resolution (compared
to orbital spacecraft observations, for example), each pixel is
unlikely to contain spectrally pure but spatially unresolved sur-
face material (with the possible exception of relatively large
expanses of CO2 and H2 O ices at the polar caps). The surface of
Mars is made up of many different minerals of differing com-
position, grain size, and crystalline structure, so surface spectra
should not necessarily be expected to precisely match labora-
tory references of pure samples. Nonetheless, it has been known
for decades that the overall average nature of the non-ice sur-
face spectrum of Mars is similar to the spectra of many kinds
of ferric–iron bearing minerals (see review above). Spectral re-
flectances typical of these minerals rise smoothly through the
ultraviolet into the red regions (600–700 nm), though the slope
varies from region to region on Mars and particularly between
the sharply demarcated bright and dark portions of the surface.
Further, most ferric oxides are characterized by a reflectance
peak between 700 and 800 nm (e.g., Morris et al., 1985, 2000).
The challenge in interpreting these and similar observations is
to identify specific surface constituents through discernable and
diagnostic spectral features, and to map areas where these con-
stituents are relatively more prevalent, despite the mixing of
several 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 area
Fig. 9. (Top) The spectra of the bright area in the vicinity of Olympus/Amazonis in each pixel is measuring the band depth. In our band depth
acquired by ISM (upper curve) in 1989 and STIS (lower curve) in 1999. The calculations, a straight-line continuum is drawn between two
dashed line is the ISM spectrum translated to equal the reflectance in the STIS points on a spectrum, and the difference between an observed
data at approximately 840 nm. (Center) The correction vector produced by di-
value at some point between the endpoints of the continuum
viding ISM spectrum values by STIS values at wavelengths longer than 840 nm.
(Bottom) The resulting “Bright Mars” composite spectrum used in spectral ra- and the continuum value at that point is measured (e.g., Clark
tio curves for comparison with terrestrial analogs to martian surface materials. and Roush, 1984; Bell and Crisp, 1993). This difference is
expressed as a fraction (here a percentage) of the continuum
gion spectrum from Mustard and Bell (1994) to generate a value, such that a band depth of zero is no band, positive band
single “correction vector” that could be applied to all of the depth indicates a concave (absorption) feature, and negative
spectra in the 1999 STIS data set. We used the Mustard and Bell band depth indicates a convex feature.
(1994) Olympus/Amazonis spectrum (Spot 41) for our com- Band ratios can also provide a useful tool to search for min-
parison spectrum, because we observed roughly the same high erals based on diagnostic spectral shapes, as well as a way to
enhance differences in surface reflectance in order to visually
albedo, presumably dust-covered, surface area in the 1999 STIS
identify faint surface details. In this research, we used band ra-
data. A STIS “Bright Mars” spectrum was derived by averaging
tioing to look for detail in the ultraviolet and short-wave visible
∼20 pixels centered on the Olympus/Amazonis region (Fig. 9).
spectra collected in 2001 and 2003; we divided the I /F values
The average I /F value at 839 nm for that region in the STIS
in one chosen band by those in another. A related technique was
data was 0.287, and the corresponding ISM-telescopic spec- used to search for sub-pixel surface constituents by dividing
trum’s I /F at 839 nm was 0.298, a difference of only about 4%. spectra in all pixels by the spectrum of a representative bright
To build the correction vector, the ISM-telescopic spectrum was martian region presumed to be dominated by ferric aeolian dust
scaled to the STIS average spectrum at 837 nm then divided by (cf. Merényi et al., 1996). While the absolute values of spectral
the corresponding STIS I /F values. Finally, all spectra in the ratio curves shown here have little meaning when the ratios are
1999 STIS data set were adjusted in the range 837–1027 nm between portions of the martian surface and laboratory spectra
by multiplying all of the spectra by this one constant correction of reference materials, relative differences in absorption band
vector. I /F values at wavelengths less than 837 nm were not position or spectral slope can be used to infer mineralogic sim-
modified. It is important to point out that this scaling method ilarities or differences. Thus, we scaled these curves to unity
only changes the absolute shape of I /F data longward of 837 at 529 nm to facilitate comparison between laboratory mineral
nm, but it does not change any of the relative region-to-region samples and generally less reflective martian surface measure-
spectral variations inherent in the data set. ments.590 J.F. Bell III, T.M. Ansty / Icarus 191 (2007) 581–602
We explored the possibility of using more complex analysis
methods such as principal components analysis (PCA) or linear
spectral mixture modeling on our STIS data. However, none
of these methods revealed useful new information relative to
the much simpler approaches outlined above. For example, un-
raveling the physical meaning behind the eigenvectors derived
by PCA methods was problematic for these hyperspectral data
sets. Also, it was not possible to untangle the effects of grain
size from abundance in simple linear mixing models at these
wavelengths. Thus, since our focus is primarily on relative re-
gional spectral differences, our simpler ratio, band depth, and
spectral 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 a
steady positive spectral slope across the near-UV through the
visible region to about 600 nm, followed by a reflectance
maximum 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). Spatial
variations in this slope reveal interesting features despite the
broadly-similar nature of almost all surface spectra.
Most ferric oxide/oxyhydroxide minerals possess reflectance
maxima between 700 and 800 nm (e.g., Morris et al., 1985,
2000). A reflectance peak appears as a negative value of band
depth when that parameter is defined relative to continuum val-
ues on either side of the peak. For example, Fig. 10 shows a map
of the band depth at 759 nm defined using continuum points at
Fig. 10. (a) Band depth at 759 nm, measured against a continuum between 701
701 and 818 nm, generated from our 1999 STIS data set. While and 818 nm. Bright areas are those showing the least negative band depth (low-
virtually all of Mars shows negative band depth as defined this est absolute value), and dark areas the most negative band depth. The range of
way at this wavelength, the areas most negative are those with the image is −3.4% to −0.5% band depth. This is a Mollweide equal-area map
the highest albedo. Areas exhibiting less negative 759 nm band projection of the data, with 90◦ N at the top, 90◦ S at the bottom, 0◦ longitude
depth values (brighter areas in Fig. 10) are lower albedo regions running vertically down the middle, and 180◦ longitude along the periphery
of the map. (b) Global albedo map in the same projection, from MGS/TES
that also show evidence (as described below) for increased near- (Christensen et al., 2001). (c) Global topography map in the same projection,
IR band depth attributed to ferrous iron-bearing minerals like from MGS/MOLA (Smith et al., 1999). White is high, black is low.
pyroxene.
The spectral slope between about 530 and 570 nm in gen-
eral correlates with albedo: the more reflective areas in this 833 nm it reaches 0.31, among the highest values on the planet
range show proportionally higher spectral slope than darker re- at those wavelengths. Spectral curves for the volcanoes appear
gions. Exceptions were observed in Mars’ western hemisphere, similar to but translated to slightly higher values than the curves
where Olympus Mons and the three volcanoes in the Tharsis for surrounding terrain, at least partially explaining the lower
Montes possess notably lower spectral slope than surrounding 570 to 530 nm band ratios.
regions of roughly similar albedo. For example, the ratio of Possible explanations for the higher I /F values over the vol-
I /F at 570 nm to I /F at 530 nm for Olympus Mons is ap- canoes include more reflective ice clouds near the summits,
proximately 1.35, and the value near the summit of Ascraeus increased opacity of airborne dust with lower elevation, con-
Mons (the northernmost of the three large volcanoes in Tharsis tributing to lower I /F at lower altitudes, or a layer of surface
Montes) was 1.30. These values are significantly lower than that frost across the mountaintops. As the 1999 observations oc-
in nearby high albedo terrain (∼1.47). For comparison, values curred during martian summer, the volcanoes are equatorial
from Syrtis Major and Arabia (classical low and high albedo re- features, and no typical water or CO2 ice spectra were observed,
gions, respectively) were 1.15 and 1.43, respectively. Radiance frost is an unlikely explanation. Likewise, the relative clarity of
factor values from the summit of Olympus Mons are higher than surface detail compared to observations made during periods
surrounding terrain; at 657 nm the I /F value is 0.27 and at of higher airborne dust activity (Fig. 6), and the MGS/TES-STIS Mars observations 591
derived record of low atmospheric dust opacity during this time Earth-based viewing geometry than in nadir-viewing orbital
(e.g., Smith, 2004) argue against dust obscuration of lower al- imaging.
titude terrain. Clouds are frequently observed in the vicinity of
Mars’ higher peaks, especially during the aphelion season (e.g., 3.2. Vis-NIR reflectance maxima (M1) and minima (T1)
Wolff et al., 1999) and thus thin water ice clouds are the likely
cause for these higher reflectance measurements. As described above, martian surface reflectance spectra at
A green to blue color ratio image (566/489 nm) from our STIS spatial resolution are broadly similar, and show a max-
2003 STIS observations reveals the presence of a subtle lin- imum value between 700 and 800 nm. Morris et al. (2000)
ear albedo feature in the high albedo Arabia region (near 15◦ examined the position of this reflectance maximum (termed
N, 355◦ W) that does not appear to have been seen in pre- “M1”) between 750 and 800 nm in Mars data from the Im-
vious STIS or MOC data sets to correspond to any obvious ager for Mars Pathfinder (IMP) instrument, comparing those
surface topographic or geologic feature (Fig. 11). The feature results to the M1 value for terrestrial ferric minerals and SNC
runs from northwest to southeast and does not appear to be an meteorites. They found that M1 could be the result of a com-
artifact of the data collection scheme, calibration, or the map bination of two or more superposed ferric mineral spectra, or
projection, as it is roughly perpendicular to the orientation of the interaction between a ferric absorption edge and the high-
the STIS slit and does not possess the smooth edges typical energy wing of a ferrous absorption feature. Because the 1999
of projection artifacts. Band ratio values (566/489 nm) within STIS data, like the IMP data set, do not extend significantly into
the feature average 1.805 ± 0.020, while northeast and south- the 1 µm band region, this ambiguity precludes mineral identi-
west of the feature the values were 1.859 ± 0.010 and 1.847 ± fication based solely on the location of this peak reflectance.
0.013, respectively. The feature is barely discernable in the 8- However, the STIS data, with much greater spectral resolu-
band aggregated (2 nm spectral resolution) data at 566 nm, and tion than IMP, allow more precise measurement of M1, and
the pattern disappears entirely as wavelength decreases to ∼536 the near-global coverage of the data set enables the trends to
nm. We searched for but could not find evidence of the fea- be assessed as a function of albedo or other parameters. As
ture in MOC/WA high resolution images from 1999 (Fig. 11). shown in Fig. 12, the position of M1 is linearly proportional
Nor was the feature seen in a search of MOC/WA images from to the radiance factor at M1. This relationship holds across
2003 around the time of our STIS observations (B. Cantor, all data, whether from high or low albedo regions. The pres-
personal communication, 2007). The lack of evidence for the ence of an M1 feature most likely implies the presence of at
feature in MOC images and the lack of a correlation to any least some crystalline ferric or ferrous minerals on the sur-
obvious, permanent surface geologic feature suggest that this face (Morris et al., 2000). The observed cluster of M1 values
low albedo marking may either be the result of small changes for low albedo regions (∼725 to 750 nm) is consistent with
in the surface dust distribution in Arabia (perhaps as a result the presence of ferric-bearing minerals like hematite, goethite,
of the planet-encircling dust storm of 2001), or that it may be and schwertmannite, and the cluster of M1 values for high
the result of a viewing geometry effect along a subtle albedo albedo regions (∼770 to 800 nm) is consistent with the presence
boundary that is simply more apparent in the specific late 2003 of ferric-bearing minerals like schwertmannite, lepidocrocite,
maghemite, and ferrihydrite (e.g., Fig. 4 of Morris et al., 2000).
Fig. 11. Band ratio image (566:489 nm) of Mars from the 2003 STIS data (top, Fig. 12. The wavelength of the peak I /F value (M1) plotted against the I /F
Mollweide projection). The dashed rectangle encloses the area enlarged at bot- value at that peak. Data points represent averages from 30 regions of interest en-
tom right, showing the linear albedo feature in the Arabia region (arrow, at compassing 1245 spectra from the 1999 STIS data set. Overall goodness-of-fit
15◦ N, 355◦ W). At bottom left, a Mars Global Surveyor (MGS) Mars Orbiting is 0.925. Separate regression curves for low albedo and high albedo subsets of
Camera Wide Angle red filter “geodesy campaign” mosaic of the same region, the data are also shown. While there is more scatter in the fit to these subsets,
generated from images acquired in 1999 (MSSS/JPL/NASA). the same overall correlation is still observed.592 J.F. Bell III, T.M. Ansty / Icarus 191 (2007) 581–602
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, Morris
et 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 and
terrestrial analog sample spectra. In particular, ferric mineral-
ogy can be better constrained using both the M1 reflectance
maximum and the so-called “T1” reflectance minimum, which
occurs between 750 and 1000 nm in many martian and lab-
oratory mineral spectra. We measured the wavelengths of the
M1 and T1 maxima and minima in the 1999 data by fitting Fig. 13. Distribution of the T1 reflectance minimum, with brighter values indi-
cating longer-wavelengths. The scale extends from 850 (black) to 970 (white)
spectra at each pixel to a sixth degree polynomial curve (in
nm. Black areas within the image either did not possess a minimum between
order to minimize the effects of noise in determining minima 800 and 1021 nm or the minimum could not be reliably measured. Note that T1
and maxima). M1 values were sought between 529 and 900 determination was highly sensitive to minor calibration variance; the diagonal
nm, while T1 values were sought between 800 and 1020 nm. banding in the northern hemisphere is not correlated with known or observed
The resulting distribution of M1 values is very strongly cor- surface features.
related with albedo, with classical bright regions exhibiting
higher 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 surface
in the high albedo regions being dominated by ferric minerals,
and the low albedo regions showing more evidence of ferrous
minerals.
The distribution of T1 reflectance minima values is not as
strongly correlated with albedo (Fig. 13), and also appears to
be sensitive to minor variations in observed I /F caused by cal-
ibration variance and different observation dates, resulting in
a higher level of artifacts in the distribution map. Plotting T1
against M1 as a histogram (Fig. 14) reveals a rather compact
Fig. 14. Two-dimensional histogram showing the correlation of measured M1
and correlated distribution of values. High values of M1, which
I /F maxima and T1 I /F minima in our 1999 HST/STIS Mars data set. Darker
tend to correspond to low values of T1, are found in high albedo pixels correspond to higher histogram bin values.
regions of Mars; conversely, low M1/high T1 points are from
the darker portions of the surface. The range of M1/T1 val- the Opportunity rover (e.g., Klingelhöfer et al., 2004). One of
ues plotted in Fig. 14 is consistent with the range derived for
the motivators of our STIS UV observations was that jarosite
candidate martian ferric and ferrous minerals by Morris et al.
possesses a very characteristic 430–440 nm absorption band
(2000). However, despite the near-global coverage of the data
which could be a key indicator of its presence even in a
set and the relatively high (by Earth-based standards) spatial
highly mixed pixel (Fig. 15; Table 2; Sherman et al., 1982;
resolution of the observations, our analysis does not reveal spe-
Sherman and Waite, 1985; Morris et al., 2000). We searched
cific regions of M1/T1 space that could be used to discriminate
for but could not detect any evidence of this feature in the 2003
among the specific kinds of ferric and ferrous minerals studied
by Morris et al. (2000). Rather, the correlated and somewhat STIS data set using band ratios, spectral slope measurements,
bimodal nature of the M1/T1 distribution simply provide addi- and band depth mapping.
tional support for the ferric-dominated nature of bright region Although the jarosite feature near 435 nm is unique and di-
Vis-NIR reflectances and the likelihood that ferrous minerals agnostic among ferric oxides, noise in the 2003 data set in the
are important Vis-NIR spectral components of the darker, pre- range between 400 and 440 nm is a significant obstacle to ob-
sumably bedrock areas. serving this feature. The average I /F of the martian surface
rises from about 0.04 to 0.06 over that range (e.g., Fig. 24).
3.3. Searching for jarosite Reference spectra of pure jarosite samples show a peak re-
flectance of 0.14, surrounded by shoulders with approximately
As noted above, jarosite is a ferric–iron bearing hydrated sul- 0.10 reflectance, for a peak magnitude of 1.4 times the base re-
fate of particular interest on Mars, as its presence had been flectance (Fig. 15). Averaging I /F spectra across all of Mars
predicted based on some Mars chemical weathering models within this data set (8-band aggregate data, spectral resolution
(e.g., Burns, 1987) and ultimately it was discovered to make approximately 2 nm), the magnitude of fixed I /F noise arti-
up a small fraction (STIS Mars observations 593
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: Narrow
Fe2+ 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, Thus, specific evidence of jarosite is not present in our 1999
individual pixels exhibit random noise fluctuations on the or- or 2003 UV to near-IR STIS data sets, either based on the
der of 0.001 I /F . These noise sources result in an average 435 nm or the 705 nm reflectance features. Seeking evidence
signal-to-noise ratio of 40 to 60 (based on random noise) or 4 for jarosite in our airborne dust-dominated 2001 STIS UV data
to 6 (considering the systematic noise artifacts in the corrected set was no more fruitful. Band depth, band ratio, and spectral
data). Expressed another way, the random noise magnitude is slope measurements generated from that data set in the appro-
approximately 4% of a baseline reflectance value of 0.05. As priate range showed only map-projection artifacts and low SNR
the jarosite peak is 40% greater than the baseline reflectance, artifacts.
jarosite abundance at the 10% level over broad areas of the mar-
tian surface (comparable to the resolution of our data) would be 3.4. Searching for pyroxene
necessary to be discerned at the same level as random noise
in the data; much higher abundance levels would need to be As discussed above, the strength of the near-IR ferrous ab-
sorption feature (“1 micron band”) can be used to assess the
present to exceed systematic noise artifact levels within our data
ferrous iron-bearing mineralogy of the surface materials. For
set. These abundances are not indicated either by in situ rover
example, in our 1999 Vis-NIR STIS data, band depth at 900 nm,
measurements nor other remote observations of the planet. Fur-
measured against a continuum between 741 and 1021 nm, cor-
thermore, band depth measurements across the jarosite feature
relates very strongly with the classical low albedo areas of Mars
would show some spatially coherent patterns if the jarosite frac-
(Fig. 16). Our STIS 900 nm band depth map exhibits many sim-
tion were at approximately the same level as random noise; no
ilarities to the MEx/OMEGA equatorial to mid-latitude maps of
such spatial patterns were observed. pyroxene derived by Bibring et al. (2005, 2006).
Jarosite also exhibits a reflectance peak at 705 nm that could We generated similar band depth maps at other wavelengths
be an indicator of its presence (e.g., Morris et al., 2000). Un- between 900 and 1000 nm, searching for evidence of slight spa-
fortunately, several other ferric minerals (e.g., hematite, schw- tial variations in band center position that could correlate with,
ertmannite, and akaganeite) also possess reflectance peaks in for example, slight changes in pyroxene chemistry (Adams,
the vicinity of 705 nm, so a conclusive search for jarosite must 1974; Cloutis and Gaffey, 1991; Bibring et al., 2005). While
exclude these other phases. While virtually all areas of Mars im- we could detect differences in the absolute band depth values,
aged in the 1999 data set show positive band depth at 705 nm, we could not identify any statistically-significant spatial differ-
no coherent spatial feature was found corresponding to a maxi- ences in band depth maps made at these different band cen-
mum reflectance value at this wavelength within narrow search ters. Unlike OMEGA spectra (e.g., Bibring et al., 2005, 2006;
bounds (∼20 nm). Additionally, the M1/T1 characteristics that Mustard et al., 2005), our STIS data may not have high enough
would indicate discernable quantities of jarosite did not occur SNR, or extend to long enough wavelengths, to allow us to de-
at surface locations where spectral evidence of jarosite could be tect subtle spatial variations in the martian spectrum in the 900
found. to 1000 nm region. Thus, we use 900 nm band depth (Fig. 16)594 J.F. Bell III, T.M. Ansty / Icarus 191 (2007) 581–602
Fig. 16. A band depth image of Mars, measuring depth at 900 nm as a fraction of
the continuum value between 740 and 1021 nm. Areas in black 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 the
largest 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. Low
albedo northern hemisphere regions (besides Syrtis Major)
show structured high 900 nm band depth as well, but maximum
band depths are significantly lower than in Syrtis and southern
hemisphere dark regions: about 8.4%. A circumpolar ring of
low albedo material at high northern latitudes possesses a high
900 nm band depth (as high as 11%) relative to surrounding
terrains. 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 this
region in previous HST multispectral imaging studies (e.g., Bell
et al., 1997a, 1997b). In contrast, high albedo regions show 900 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 relevant
nm band depth on the order of only 1–2%. The north polar cap to a search for pyroxenes (see text). At bottom (a), an averaged spectrum from
showed very slight negative band depth, presumably a mani- areas with greater than 5.5% band depth at 900 nm, measured against a con-
festation of ices, rather than basaltic minerals, dominating that tinuum between 740 and 1021 nm, which do not show increasing reflectance
region’s near-IR spectral characteristics. Representative exam- between 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,
ples of these spectra are shown in Fig. 17.
but with the more characteristic inflection near 900 nm. Spectrum (c) is from
Spectra extracted from pixels showing the highest 900 nm the north polar sand sea (22 pixel area, location approximately 75◦ N, 45◦ W),
band depth (Fig. 18, lower panel) show some general similari- closely resembling the high band depth areas that were observed mostly in the
ties to laboratory spectra of pyroxenes. I /F values in the high southern hemisphere (see Fig. 16). Curve (d) is the reference “Bright Mars”
900 nm band depth spectra gradually rise through the near-UV spectrum, measured at Olympus Mons. At top, (e) is a typical spectral curve
from the north polar cap (26 pixels, at approximately 80◦ N, 300◦ W). Spec-
and visible regions to a peak value of 11–13% near 750 nm, tra b, c and d are each offset by 0.05 from the curve below them; e is offset by
show a broad I /F minimum centered near 900 nm, then begin 0.20 from d. Representative 1σ error bars are shown at every ≈50 nm intervals.
to increase again towards 1000 nm. These uncertainties were derived by normalizing all spectra from each area at
In an effort to identify whether pyroxene is indeed present 904 nm and then calculating the standard deviation of the normalized set.
in these pixels, high 900 nm band depth spectra were divided
by the composite “Bright Mars” spectrum (Fig. 9) created when near that wavelength (e.g., Morris et al., 1995, 1997, 2000). The
bootstrapping the STIS data to the calibrated ISM observations. mixed-pixel property of the STIS data, combined with Mars’
The resulting ratioed dark/bright spectra resembled neither lab- ever-present aeolian dust, assure that a ferric reflectance maxi-
oratory spectra of pure pyroxene nor that of any specific ferrous mum in the vicinity of 750–800 nm will influence spectra every-
(or ferric) mineral (Fig. 18, upper panel). The presence of a where at this spatial resolution.
minimum value near 950 nm in some of these ratio spectra is Burns (1993) noted the potential of 506 and 548 nm ab-
consistent with pyroxene, but could also be consistent with an- sorption features as secondary identification features of Fe2+ in
other mineral or mixture of minerals with an absorption feature orthopyroxene laboratory spectra (Fig. 15). Unfortunately, nei-You can also read
