esa mars research abstracts part 1
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
3D HOT PARTICLE AND UPPER ATMOSPHERE MODELLING OF MARS H. I. M. Lichtenegger1,
H. Lammer2, H. Gröller
3, Yu. N. Kulikov
4,1Space Research Institute, Austrian Academy of Sciences,
Schmiedlstr. 6, Graz, Austria,2Institute for Geophysics, Astrophysics and Meteorology, University of Graz,
Austria,3Polar Geophysical Institute, Russian Academy of Sciences, Khalturina Str. 15, 183010, Murmansk,
Russian Federation, [email protected]
A new 3D hot particle Monte Carlo code which can
be coupled to a 3-D exosphere test particle model is
presented. These coupled codes can be used for
studying expected asymmetries related to latitude
and longitude as well as day and nightside
production rates and distributions of hot particles in
planetary exospheres. The newly photochemically
generated energetic neutral atoms are traced from
their point of origin up to the exobase as a function
of longitude, latitude, production process, collision
probability with the cool background atmosphere,
change of direction (altitude and angles) andenergy dependent collision cross sections. For
modelling the Martian background atmospheric and
temperature profiles from the mesopause to the
exobase we apply a diffusive gravitational
equilibrium and thermal balance model. The hot
particles which arrive above the exobase with
energies higher than the corresponding exobase
temperature of the background gas are divided into
energy bins and used for the calculation of the
energy density distributions as a function of
latitude and longitude. These calculated energy
density distributions of photochemically produced
hot atoms at the Martian exobase are used as inputsfor 3-D hot particle exosphere simulations. Finally
we compare our results with that obtained with two
stream models.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ALLUVIAL FAN AND DELTA PROGRADATION IN MARTIAN CRATER LAKES
M. G. Kleinhans1, H. E. Van de Kasteele
2, E. Kraal
3.
1Universiteit Utrecht. Faculty of Geosciences, PObox
80115, 3508 TC Utrecht, The Netherlands.2No affiliation, Odijk, The Netherlands.
3Virginia Tech Geoscience
Department, Blackburg, VA 24060, USA. [email protected]
Setting and problem: Numerous sedimentaryfans and fan deltas have been found on Mars
1,
indicating the presence of flowing water on the
surface. In principle the morphology and
dimensions of the fans and deltas are strongly
determined by the formative boundary conditions,
and not so much by the details of the sediment
transport that formed them. Upstream boundary
conditions are the flow discharge and sediment
input2
and the downstream boundary condition is the
lake water level (or the lack thereof which results in
a fan rather than a fan delta). The lake water level
depends on the water input, the crater dimensions
and infiltration, evaporation or overflow. So far theoverall volumes of fans and deltas have been used to
infer flow discharge and flow duration2. We aim to
infer combinations of these boundary conditions
from detailed morphology and report preliminary
results with an analytical method.
Analytical method: We describe idealised
sedimentary bodies by a cone, the subaerial fan with
a given gradient of, say, 0.05, on top of a
horizontally truncated cone, the subaqueous delta
with a given gradient of, say, 0.1-0.6 (up to the
angle of repose of the sediment) (Figure 1a). By
these volumes we derive an analytical cubic
equation for the volume of the idealised delta or fan(valid for both). The necessary input parameters are
the fan surface gradient (from observations or
dependent on sediment transport capacity) and the
clinoform gradient (from observations or at most the
angle of repose of unconsolidated noncohesive
sediment). The lake level determines the elevation
of the break in gradient (the shoreline). The
shoreline position is determined by the sediment
input (given or calculated2), its duration, and the
lake level.
The lake level is calculated from the inputdischarge and the crater dimensions. The crater
volume can be calculated from empirical power
functions for crater depth (from diameter) and cross-
sectional profile3. We integrated the profiles to
obtain the volume of the lake. Given a water input,
the lake level rises at a variable rate over time
(Figure 1b). A unique shape, shoreline position and
height of a fan or delta is now calculated by the
standard solution (first root) of the cubic equation
(Figure 1c). If this solution is negative then there is
not enough sediment to form a subaerial part and the
system is drowned and deposited as a simple cone.
Results: The calculations are presented for theexample of figure 18d in ref.1, which is a lobed fan
delta in a crater of diameter D=64km and depth
d=1.9km. We reconstructed a flow discharge of
250000 m3 /s and a sediment input of 0.011 km
3 /day
(reported elsewhere4). The rate of water level rise
declines because the crater widens (figure 1b).
When the water level is still low, the delta is long
but not high, but while the water level rises the
sedimentation cannot keep up so that the shoreline
retreats, the delta height increases but its length
diminishes (figure 1c). The end result for this and
many other reasonable combinations of water and
sediment input is a steep fan wherein the clinoformsare buried. This result suggests that some fans
observed in Martian craters may have formed as
drowned deltas. The most important uncertainty is
how much water is lost through infiltration,
evaporation and through the crater rim.
References: [1] Irwin, R.P., A.D. Howard, R.A.
Craddock, and J.M. Moore (2005), JGR 110, E12S15,
doi:10.1029 /2005JE0024 60. [2] Kleinhans, M.G.
(2005), JGR 110, E12003, doi:10.1029/2005JE002521.
[3] Garvin, J.B. and J.J. Frawley (1998), GRL 25, 24,
4405-4408. [4] Kraal, E. et al. AGU fall meeting 2007
Figure 1a. Fan delta described by top cone on truncated cone. Shoreline position (circle) depends on water and sediment
input and on lake level history, while the latter depends on water input and crater dimensions.
b. Water level rise for a crater of diameter D=64km, depth d=1.9km for a flow discharge of 250000m3
/s.c. Resulting delta evolution in a simple basin: the delta drowns (same vertical scale as b) and the shoreline (circles) retreats.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
The alteration of the mineralogical composition of ejecta at Syrtis Major D. Baratoux1, P.C. Pinet1, A.
Gendrin2
L. Kanner2
J. Mustard2
Y. Daydou1
J. Vaucher1
and J.-P. Bibring3
1Observatoire Midi-Pyrénées, UMR 5562 CNRS, Laboratoire Dynamique Terrestre et Planétaire, Université
Paul Sabatier Toulouse III, Toulouse, France. 2Department of Geological Sciences, Brown University,
Providence, Rhode Island, USA. 3Institut d’Astrophysique Spatiale, UMR 8617, Orsay, France.
We summarize in this paper the recent
observations which have been made from OMEGA
[1] and THEMIS-IR data [2]. A large number of
impact craters at Syrtis show a thermal and spectral
signature which seems to be weakly affected by the
erosional history.
Figure 1. RGB pyroxene map from OMEGA data. Blue,
high-calcium pyroxene band strength; red and green, low-calcium pyroxene band strength. For some of the largest
craters, the ages have been derived from crater counts and
are indicated. Ages for type I ejecta (HCP-rich) appear in
blue, while ages for type II ejecta appear in orange. The
full-resolution OMEGA has been represented on aTHEMIS-IR daytime image for the crater 0738 + 114
(RGB composition with a different stretch, data from orbit
number 444)
A RGB-composite map showing the High-
Calcium-Pyroxene abundance versus Low-Calcium-
Pyroxene is presented on the Figure 1 showing the
enrichment of some impact ejecta relative to theSyrtis lava flows. In particular, the OMEGA
spectrum and inferred composition display an
axisymmetric pattern which can be explained by the
excavation in a terrain having a variable
composition with depth followed by the ejecta
emplacement flow (Figure 1).
Then, it is shown that the enrichment in High-
Calcium pyroxene corresponds to the younger ejecta
(Figure 2). Several hypotheses are discussed
concerning the reasons why the spectral signature of
this enrichment is now masked to the observations
from the orbit, including the presence of dust, and
possible processes of alteration in the present cold
and dry environment.
Figure 2. Ages of HCP-rich craters versus other craters.
References: [1] Baratoux, D., N. Mangold, P. Pinet, and
F. Costard (2005), Thermal properties of lobate ejecta in
Syrtis Major, Mars: Implications for the mechanisms of
formation, J. Geophys. Res., 110, E04011,doi:10.1029/2004JE002314. [2] Baratoux, D., P. Pinet, A.
Gendrin, L. Kanner, J. Mustard, Y. Daydou, J. Vaucher,
and J.-P. Bibring (2007), Mineralogical structure of the
subsurface of Syrtis Major from OMEGA observations of lobate ejecta blankets, J. Geophys. Res., 112, E08S05,
doi:10.1029/2007JE002890.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ANALYSIS OF CO2 NON-LTE EMISSION AT 4.3 μm IN THE ATMOSPHERE OF MARS WITH PFS
LIMB OBSERVATIONS V. Formisano1, M. Grossi
1, M.A. Lopez-Valverde
2,G. Gilli
2, M. Giuranna
1.1
IFSI-
INAF, via del Fosso del Cavaliere 100, 00133, Rome, Italy.2
Instituto de Astrofisica de Andalucıa (CSIC),
Apdo. 3004, Granada, Spain. [email protected]
We present PFS-MEX limb observations of the CO2 non-local thermodynamic equilibrium (non-
LTE) emission at 4.3 μm in the atmosphere of Mars
collected in more than one Martian year (see
Formisano et al 2006). The spectral shape along
each orbit changes with the altitude of the tangent
point, giving indication on how the major and minor
CO2 bands and/or isotopes contribute to the
emission at different heights. We analyze how the
solar zenith angle (SZA) affect the intensity of the
observed radiance and their contributions.
There is evidence for a clear dependence of the
non-LTE peak emission on the SZA, showing a
cosine-like relation. Moreover, the altitude of themaximum emission along the orbits, approximately
constant at 100 km altitude up to 70o
SZA, lowers
for larger SZA, decreasing even down to 50 km at
SZA = 88o
in some orbits. We discuss how such
observations can be interpreted in terms of the non-LTE theory.
These results, while on one hand confirm and
quantify some aspects of the non LTE computation,
on the other hand will further stimulate theoretical
modelling , possibly bringing closer the moment in
which the measurements could be inverted to obtain
important information about the high altitude
atmospheric properties.
References: Formisano , V., Maturilli A. , Giuranna M.,
D’Aversa E., Lopez-Valverde M.A. . “ Observations of
non-LTE emission at 4-5 microns with the planetariFourier Spectrometer aboard Mars Express mission.
(2006), Icarus 182, p. 51-67.
Figure 1. Radiance of the peak emission at 2340 cm1
(right panel) and at 2316 cm1
(left panel) versus thecorresponding solar zenith angle. There appears to be a clear dependence of the radiance on the cosine of theSZA. The dashed lines give the overall fit to the data. The fit parameters are shown in the left-bottom corner of each plot.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ANALYSIS OF SNOW ALGAE FROM THE CAIRNGORM MOUNTAINS, SCOTLAND: HIGH
RESOLUTION MEASUREMENT OF ORAGNICS IN ICY TERRAINS S. J. M. Phillips & J. Parnell.
Department of Geology and Petroleum Geology, School of Geosciences, College of Physical Sciences, Meston
Building, King’s College, Aberdeen, AB24 3UE, Scotland, UK. [email protected]
Introduction:Snow algae are extremophiles that live and grow
in semi-permanent to permanent snow or ice in the
alpine or polar regions of the world and are also
found in old snow beds on the Cairngorm plateau,
Scotland. Hence, the Martian ice-caps are a potential
niche for snow algae. Their optimum growth
temperatures are below 10ºC. To survive in this
extreme environment they require several
adaptations. These include the biosynthesis of
pigments, polyols, sugars and lipids, mucilage
sheaths, motile stages and spore formation [1].
Snow beds are a good target for the detection of
life since snow algae occur with large cellabundances. Cell concentrations of 10
5to
10
6 per ml
are common in large blooms.
This study is aiming to investigate the sites, the
concentration and the character of snow algae in the
semi-permanent Cairngorm snow beds. The snow
algae used in this study is also to test a range of
analytical techniques that could potentially be used
in planetary exploration.
The Cairngorms in Scotland:
The Cairngorm Mountains in north east Scotland
is the only sub-arctic site in the UK, providing a
unique climate, ecology and terrain, with a mixtureof continental and oceanic climates. The area
receives over 100 days of snowfall per year and has
the only perennial snow cover in the UK [2].
Method:
Red snow algae (Fig. 1) located on Cairn Gorm
in the Cairngorm Mountains was aseptically
collected in July 2007. The sampling site was
located 150 m north east from the summit of Cairn
Gorm at an altitude of 1100 m above sea level. The
site was in a sheltered aspect, protected from the
prevailing wind, where snow from the previous
winter had survived through until late summer. Thesamples were then transferred in a frozen state to the
laboratory and stored at -10ºC. The samples have
been analysed for biological markers with SERS
(Surface–Enhanced Raman Spectroscopy) and GC-
MS (Gas Chromatography–Mass Spectrometry).
Figure 1. Red snow algae, Cairn Gorm, Scotland, UK.
Results:
Field observations of the snow algae in situ
reveal that the algae are red to brown in colour and
occur as irregular patches of approximately 1 m in
diameter. The algae exist on the snow surface and
extend down below the surface of the snow to a
depth of 3 mm, which is within the photic zone.
SEM observation of the algae (Fig. 2) reveal that the
algae are round in shape and range in size between 5
and 16 μm. The samples are being analysed further
at the University of Aberdeen.
Figure 2. SEM (Scanning Electron Microscope) image of
cryophilic Chlorophyceae (cC) from a snow bed on Cairn
Gorm, Scotalnd, UK.
References: [1] ezanka, R., Nedbalová, L. & Sigler, K.,
Microbiol Res 2006; (2007),
doi:10.1016/j.micres.2006.11.021. [2] Phillips, S.J.M. &
Parnell, J. (2006), LPSC XXXVII , Abs. #1027.
20 mm
cC
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ANCIENT HEAT FLOW AND CRUSTAL THICKNESS AT THE AMENTHES REGION, MARS. Javier Ruiz1, Carlos Fernández2, David Gomez-Ortiz3, James M. Dohm4,5, Valle López6, Rosa Tejero7 1Museo Nacional de Ciencias Naturales, CSIC, 28006 Madrid, Spain. 2Departamento de Geodinámica yPaleontología, Universidad de Huelva, 21071 Huelva, Spain. 3ESCET-Área de Geología, Universidad Rey JuanCarlos, 28933 Móstoles, Spain. 4Department of Hydrology and Water Resources, University of Arizona, Tucson
85721, AZ, USA. 5Lunar and Planetary Laboratory, University of Arizona, Tucson 85721, AZ, USA.6Seminario de Ciencias Planetarias, Universidad Complutense de Madrid, 28040 Madrid, Spain. 7Departamentode Geodinámica, Universidad Complutense de Madrid, 28040 Madrid, Spain. [email protected],
The Amenthes region is adequate for analyzingthe thermal structure and thickness of the Martiancrust, since estimations of both the brittle-ductiletransition depth [1,2, this work] and the effectiveelastic thickness of the lithosphere [3-5, this work]are possible for the Late Noachian/Early Hesperiantime. As such, we analyze the Late Noachian/EarlyHesperian surface heat flow of the Amenthes region
by considering homogeneously distributed crustalheat sources (and linear thermal gradients for theupper mantle), which have abundances based in thelatest GRS data reported in [6], and crustal andlithospheric mantle contributions to the totalstrength, and hence to the effective elastic thickness,of the lithosphere [7,8]. This permits us to constrainthe thickness of the Martian crust in a wayindependent from previous works. We also considerdry and wet rheologies for the lithosperic mantle.
The depth to the brittle-ductile transition deducedfrom modeling of the topography of AmenthesRupes is 25-40 km (with values of ~25-30 km being
the most probable), and the associated surface heatflow is 26-36 mW m-2 (for a crustal thermalconductivity of 2 W m-1 K-1). On the other hand, theeffective elastic thickness in this region is 19-33 km:the surface heat flow deduced by considering crustaland lithospheric mantle contributions to the totallithospheric strength, as well as wet or dry olivinefor lithospheric mantle rheology (and a lithosphericmantle thermal conductivity of 3.5 W m-1 K-1), is34-45 mW m-2 .
It is clear the narrow range of values for whichthe heat flow obtained for the Amenthes region fromthe effective elastic thickness is consistent with thatdeduced from the depth to the brittle-ductiletransition. By taking simultaneously into accountcalculations based on both metodologies, a surfaceheat flow of 35-36 mW m-2 (with a high fractionoriginated from crustal heat sources), a wet mantle
rheology, and a local crustal thickness is 45-60 kmare obtained.
A wet lithospheric mantle rheology is consistentwith results of comparisons of effective elasticthickness evolution through time with thermalhistory models for Mars [9,10]. On the other hand,our results suggest an average thickness of ~40-60km for the Martian crust (the thickness of the crust
in this region is ~0-5 km thicker than the averageplanetary value [11]), which is consistent with therange of 38-62 km obtained for [12] fromsimultaneously considering several geophysical andgeochemical arguments.
The obtained mantle heat flow, ~4-9 mW m-2, islow compared with the predictions from mantleconvection models for Mars [13], which could be alocal (and maybe temporal) phenomenon.Alternatively, the emplacement of a substantialfraction of radioactive heat sources in the crustcould have contributing to the slugging of mantleconvection [14].
References: [1] Schultz, R.A. and Watters, T.R.(2001), GRL 28, 4659-4662. [2] Grott, M. et al. (2007), Icarus 186 , 517-526. [3] Watters, T.R. (2003), Geology
31, 271-274. [4] Watters, T.R. and McGovern, P.J.(2006), GRL 33, 10.1029/2005GL024325. [5] Milbury,C.A.E. et al. (2007) PSS 55, 280-288. [6] Taylor, G.J. etal. (2006), JGR 111, 10.1029/2005JE002645. [7]Ruiz, J. et al. (2006), Icarus 180, 308-313. [8] Ruiz, J. etal. (2006), J. Geodyn. 41, 500-509. [9] Guest, A. andSmrekar, S. (2007), PEPI , in press. [10] Grott, M.and Breuer, D, (2007), Icarus, in press. [11]Neumann, G.A. et al. (2004), JGR 109, 10.1029/
2004JE002262. [12] Wieczorek, M.A. and Zuber,M.T. (2004) JGR 109, 10.1029/2003JE002153. [13]Hauck, S.A. and Phillips, R.J. (2002), JGR 107 ,10.1029/2001JE001801. [14] Reese, C.C. et al.(1998), JGR 103, 13,643-13,657.
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EVIDENCE FOR ANCIENT MARTIAN LIFE. E. K. Gibson Jr., F. Westall, D. S. McKay, K. Thomas-Keprta, S.
Wentworth, and C. S. Romanek, Mail Code SN2, NASA Johnson Space Center, Houston TX 77058, USA.
Three SNC meteorites ranging in age from 4.5
Ga. to 1.3 Ga. to 165 m.y. contain features suggestive of
past biogenic activity on Mars [1,2]. Because we do not
know what past martian life looks like or its physical orchemical properties, the only tools or criteria which the
scientific community have to evaluate evidence of past life
is to use evidence for early life on earth. There are
features within ALH84001’s carbonate globules and the
preterrestrial aqueous alteration phases of Nakhla and
Shergotty which have been interpreted as possible
evidence for past life on early Mars [1,2].
Criteria for Past life
Over the past few decades eight criteria have
been established for the recognition of past life within
terrestrial geologic samples [3,4]. Those criteria are: (a)
Is the geologic context of the sample compatible with past
life; (b) Is the age of the sample and its stratigraphic
location compatible with possible life; (c) Does the
sample contain evidence of cellular morphology and (d)
colonies; (e) Is there any evidence of biominerals showing
chemical or mineral disequilibria; (f) Is there any
evidence of stable isotope patterns unique to biology; (g)
Are there any organic biomarkers present; (h) Are the
features indigenous to the sample? For general
acceptance of past life in a geologic sample, essentially
most or all of these criteria must be met.
ALH84001, Nakhla and Shergotty vs. the Criteria for
Past Life
How does the scientific information fromALH84001, Nakhla and Shergotty compare to the
established criteria?
Geologic context. A martian origin for the three
meteorites has been shown by their O-isotopic
compositions [5] and trapped martian atmospheric gases
[6,7]. The exact martian provenances for these igneous
rocks is unknown. However, because of its 4.5 Ga. age,
ALH84001 probably originates from the early martian
crust (i.e from the ancient southern highlands). Nakhla
and Shergotty are undoubtly from younger volcanic
provenances. The presence of secondary globules or
pancake carbonates in ALH84001 and clays in Nakhla
have been interpreted as an indication of relatively low-temperature secondary mineralization by a fluid, possibly
water [5,8]. Formation of the secondary carbonates and
preterrestrial aqueous alteration at low-temperatures from
aqueous fluids would be compatible with past life, but
would not require it.
Ages and histories. The crystallization age of ALH84001
is 4.5 Ga and the rock believed to be a sample of the
original martian crust. The sample underwent extensive
shocking around the 3.9-4.0 Ga [7,9]. Carbonate
formation occurred around the 3.94 Ga [10], shortly after
the period of extensive bombardment and during a period
when the planet had abundant water [11], greater
concentrations of atmospheric gases, and highertemperatures. This corresponds to the time when life
appeared and developed on Earth [3,12]. Evaporation of
the fluids percolating through the impact-cracked surface
could have resulted in the formation of carbonates [11-13].
The sample was ejected from Mars ~17 m.y. ago and spent
11,000 years in or on the Antarctic ice sheets. We suggest
that the geologic history of ALH84001 can be compared
with terrestrial rocks of the same age and that similar
biological processes may have been operating concurrently
on Mars and Earth. Nakhla’s age is 1.3 Ga and
Shergotty’s age is between 300 and 165 m.y. Both show
evidence of preterrestrial aqueous alteration at some
period in their history (14,15).
Cellular morphologies. Some structures resembling the
mineralized casts of modern terrestrial bacteria and their
appendages (fibrils) or by-products (extracellular
polymeric substances, EPS)[16-18] occur in the rims of
ALH84001 carbonate globules and preterrestrial aqueous
alteration regions. Other bacteriomorphs are very small
but some are within the size limit of known nanobacteria
(i.e. 100-200 nm, [19-21]). Cellular-like features as large
as 1 to 2 microns are found in Nakhla [2]. Some of the
features in ALH84001 (e.g., filaments) and Nakhla
(cellular-like) are similar to terrestrial bacteria and fossil
bacteria [3,20,22]. We conclude that the evidence forfossilized microbes and their products is not conclusive,
but cannot be readily explained by nonbiological processes
and should not be ignored.
Microbial colonies. We have proposed that some of the
features in ALH84001, Nakhla and Shergotty may be the
remains of biofilms and their associated microbial
communities [16,17]. Biofilms provide major evidence for
bacterial colonies in ancient Earth rocks [22]. It is
possible that some of the clusters of microfossil-like
features might be colonies although that interpretation
depends on whether the individiual features are truly
fossilized microbes.
Biominerals. Carbonates in ALH84001 contain apopulation of magnetites having a highly restricted single-
domain size distribution and unusual morphology
(rectangular prism) that are indistinguishable from some
known, microbially-produced terrestrial magnetites, but
match no known nonbiologic magnetite [18,19,21]. We
suggest these magnetites may be formed by biogenic
processes [21]. Other magnetite grains may inorganic
[19,21]. Whisker-like magnetites (<5% total magnetites
in carbonate) described by [23-25] may have had an origin
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EVIDENCE FOR ANCIENT MARTIAN LIFE: E. K. Gibson et al.
unrelated to the rectangular prisms [21]. Work is in
progress on searching for magnetites within alteration
phases of Nakhla and Shergotty. The recent discovery of
chains of magnetites on the surfaces of carbonate globules
[26], which resemble the magnetosome chains of
magnetotactic bacteria, provide additional support for
biogenic activity within ALH84001. The discovery of
single domain, chemically pure magnetite, within
carbonate globules known to have been formed on Mars is
our strongest evidence in support of ancient martian life
[21].
Biologic isotopic signatures. Stable isotope patterns
have shown the presence of indigenous C components with
isotopic signatures of -13 to -18 ‰ [27,28], which are in
the direction of known biogenic C signatures [3].
Additional detailed study of the C-isotopic signatures is
needed to distinguish between indigenous C components
within ALH84001, Nakhla and Shergotty. Overall, the C-
isotopic signatures of the identifiable nonterrestrial,
reduced C, are compatible with biologic C-isotopic
fractionation, when compared with the signature of the
martian cabonates, but they do not prove that it occurred.
Organic biomarkers. Possible organic biomarkers are
present within ALH84001 and Nakhla in the form of
PAHs associated with carbonate globules[29]and
preterrestrial aqueous alteration regions-some of which
may be a unique product of bacterial decay [29]. PAHs in
ALH84001 are distributed in regions containing carbonate
globules [30,31] and are most likely indigenous, whereas
the other organics, such as amino acids [32] are most
likely from Antarctic contamination. Exhaustive data
must be collected before either component can be used as
a biomarker for a specific sample [33].Indigenous features. Recent studies have shown
conclusively that the PAHs are indigenous to ALH84001
and Nakhla and are not contaminants [29,31]. Based on
isotopic compositions [27,28,34] and textures, there is no
question or disagreement that the carbonate globules or
embedded magnetites in ALH84001 and the preterrestrial
aqueous alteration products in Nakhla and Shergotty were
formed on Mars and are indigenous to the meteorites.
Possible microfossil structures and some reduced C
components that are embedded in the carbonates and
preterrestrial aqueous alteration products are, therefore,
almost certainly indigenous, but other possible evidence
for life (e.g. amino acids, [32]) may be a result of terrestrial contamination [35].
Summary
Although the data are compelling, we have not
satisfied all of the eight criteria for past life described
above. However continued investiagations are in progress
and more data are needed.
Therefore, the jury is still out on early Mars life as
revealed by these meteorites [36].
We are reminded that the concept of plate
tectonics operating on the earth required 40 to 50 years
before it was accepted in the scientific community. More
recently, the hypothesis that the K-T boundary was
produced by a large bolide or comet impacting the earth
only reached acceptance after 15 to 18 years. Science does
not move swiftly in accepting radical ideas. Our
hypothesis was presented in August 1996. We believe
that after 3 years it stands stronger today than when
originally presented. To date, no fatal strikes have been
made to any of our original four lines of evidence [1],
despite several misconstrued press releases. While details
of the hypothesis are evolving as new data is generated,
we believe that our basic premise remains intact: these
meteorites contains evidence suggestive of early life on
Mars [36].
REFERENCES: [1] D.S. McKay et al., Sci. 273, 924-930
(1996). [2] D.S. McKay et al. LPSC XXX, Abst. #1816
(1999). [3] J. W. Schopf and M. Walker, In Earth's
Earliest Biosphere: Its Origin and Evolution, Ed. J.W.
Schopf, 214-239 Princeton Press (1983). [4] P. Cloud and
K. Morrison Precamb. Res. 9, 81-91 (1979). [5] C.S.
Romanet et al., MAPS 33, 775-784 (1998). [6] D. Bogard
and P. Johnson, Sci. 221, 651-655 (1983). [7] D. Bogard
and D. Garrison, MAPS 33, A19 (1998). [8] J. Valley et
al., Sci. 275, 1633-1638 (1997). [9] R.D. Ash et al.,
Nature 380, 57-59 (1996). [10] L. Borg et al., Workshop
Martian Meteorites 5-6 (1998). [11] P. Warren, JGR
103,E7,16759-16773 (1998). [12] S. L. Mojzsis et al.,
Nature 384, 55-59 (1996). [13] J. Head et al., MAPS 33,
A66 (1998). [14] J.L.Gooding et al., Meteoritics 26, 135-
143 (1991). [15] J.H.Jones, Proc. 19th LPSC 465-474
(1989). [16] D.S. McKay et al., LPSC XXVIII, 919-920(1997). [17] D.S.McKay et al. Proc. SPIE (1997). [18]
K.Thomas-Keprta et al., LPSC XXIX Abst. #1494, LPI
(CD-ROM) (1998). [19] K. Thomas-Keprta et al., LPSC
XXVIII 1433-1434 (1997), [20] K. Thomas-Keprta et al.
Geol. 26, 1031-1034 (1998). [21] K. Thomas-Keprta et
al., LPSC XXX, Abst. #1856 (1999). [22] F. Westall,
Proc. SPIE 3441, 225-233 (1998). [23] J.P. Bradley et al.
GCA 60, 5149-5155 (1996). [24] J. P. Bradley et al.
MAPS 32, A20 (1997). [25] J. Bradley et al., MAPS 33
765-773 (1998). [26] E.I.Friedmann et al., Workshop
Martian Meteorites 14-16 (1998). [27] M. Grady et al.,
Meteoritics 29, 469 (1994). [28] T. Jull et al., Sci. 279,
366-369 (1998). [29] S. Clemett et al., Faraday Disc. 109417-436 (1998). [30] G. Flynn et al., MAPS 33, A50-A51
(1998). [31] G. Flynn et al. LPSC XXX, Abst. 1087
(1999). [32] G. Bada et al., Sci. 279, 362-365 (1998).
[33] E.K. Gibson et al., Bioastr. News 10, 1-6 (1998).
[34] C.S. Romanek et al., Nature 372, 655-657 (1994).
[35] A. Steele et al., LPSC XXX, Abst. #1321 (1999). [36]
E.K. Gibson et al., Sci. Am. 277, 58-65 (1997).
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ARAM CHAOS: FRACTURING AND FLUID ACTIVITY. T.E. Zegers12
, J.H.P. Oosthoek 3, A. Rossi
2, B.
Foing2, and the HRSC Co-Investigator Team,
1Utrecht University, Faculty of Geosciences, Postbus 80021, 3508 TA Utrecht, The Netherlands,
[email protected],2ESA/ESTEC, Research and Scientific Support Department, Keplerlaan 1, 2200
AG, Noordwijk, The Netherlands, [email protected] ,3
TNO, Utrecht, The netherlands
Introduction: The chaotic terrain Aram Chaos
(2.5ºN and 338.5ºE) was mapped using MOLA,
THEMIS, MOC and in particular Mars Express
HRSC data. 3D mapping was performed using
HRSC anaglyphs and MOLA data. Special attention
was given to geometry. The area has been
previously mapped by [1]
Results: Seven units were distinguished in Aram
Chaos, with in some cases sub-units:
Highland Terrain (HT) is the main unit
surrounding Aram Chaos. Fractured HT occursupon approaching and into the Aram Chaos
depression. The fractures are 10-100 km with
fracture valley depths around 250 - 750m.
Chaotic Terrain (CT) is a lateral unit of HT. It
consists of either sharp or rounded km-scale hills
and shows the gradual loss of coherence of HT
material due to fine scale fracturing, collapse and
erosion. Although not visible in detail, the unit most
likely consists of a breccia.
Rounded Highland and Chaotic Terrain (RHCT)
is a morphological unit. It consists of rounded hills
and mesas and occurs near the Aram Chaos channel
linking Aram Chaos with Ares Vallis.
The Lower Aram Chaos Formation (LACF) is
deposited on top of the HT and consists of three
lateral units (fractured, broken and smooth). The
fractured LACF exhibits a distinct ‘glossy’
morphology visible on THEMIS VIS. It is cross-cut
by, compared to the HT fractures, relatively small-
scale fractures (1-2 km scale). Some fractures have
raised rims and some show small thrusts at the base
of the rim. The broken LACF is highly fractured,
forming ~1 km sized irregular mesas. It always
occurs at the boundary of the LACF and the
Fractured Highland Terrain. The smooth LACF isnon-fractured and may in fact be a relatively thin
unit covering the fractured LACF.
The Intermediate Aram Chaos Formation
(IACF) has a rugged morphology and is at
maximum ~250 m thick. It is stratigraphically
situated between LACF and UACF.
The Upper Aram Chaos Formation (UACF)
consists of a ~10-100 m light toned cap material,
which is relatively strong, with dark, less strong
material underneath. The cap material exhibits a
specific ‘icing’ texture and forms sharp arcuate
‘razor blade-like’ escarpments. The dark, softer
material of the UACF thickens towards the NW.
The unit is at maximum around ~300 m thick.
The Aram Ares Channel Deposits has a distinct
morphology of lineations and small elongated hills.
The Aram Chaos channel incised the AACD and
100 meter scale layering can be observed in the
channel wall.
Discussion & Conclusion: Geological mapping
using HRSC stereo image capabilities is particularly
useful to unravel the geometry of the various units
and structures in Aram Chaos as well as their cross-cutting and depositional relationships. In
combination with spectral information from
TES/THEMIS [2] and OMEGA [3], this can be used
to derive the geological evolution of the area.
The Intermediate Aram Chaos Formation
mapped here corresponds to a large extent with
areas containing hematite [2] and kieserite [3].
Concentrations of unknown hydrated minerals [3]
were found in what is mapped here as smooth
LACF. These units, as well as the overlying UACF
(Cap Unit [2]) were all deposited after fracturing
and collapse forming the chaotic terrain in Aram
Chaos. This suggests that activity of fluids, resulting
in formation of hydrated minerals in these units,
post-dated the actual chaotization process. One of
the possible mechanisms by which these deposits
could have formed is spring deposits [4].
The fracturing, brecciation, and subsidence
associated with the chaotization process are best
explained by the sudden withdrawal of water stored
in the Aram impact crater. The Aram Ares Channel
Deposit may be interpreted as the remnant of a fan
deposit formed during original flow of water into
the Aram crater. The current, eroded and channeled,
morphology of the Aram Chaos channel iscompatible with channel erosion by flow out of the
crater, probably associated with collapse and
fracturing of the overburden. Part of the trapped
fluids may have escaped by local vents around the
rim of the crater. Rapid depressurization may locally
have caused hydraulic fracturing and brecciation
resulting in the CT unit.
References: [1] Tanaka K. L. et al. (2005)
Scientific Investigations Map, 2888, USGS.
[2] Glotch T.D. and Christensen P.R. (2005) JGR,
110, E09006. [3] Gendrin A. et al. (2005) Science,
307 , 1587-1591. [4] Rossi et al., (2007),LPSC
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ASI PLANETARY MAP SERIES – MAP N° 1. GEOMORPHOLOGICAL MAP OF ARES VALLIS,
MARS. A. Pacifici. IRSPS, Università d'Annunzio, Viale Pindaro, 42, 65127 Pescara, Italy.
Introduction: Ares Vallis is one of the largest
outflow channels of Mars, which are broadcomplexes of fluid-eroded troughs. Several models
to explain the origin and the evolution of this valley
were proposed in the past, such as catastrophic
floods [1][2][3][4], glacial erosion [5][6][7], and
periglacial processes [8][9]. Nevertheless, high-
resolution images and elevation data acquired by
newest Martian missions allow studying Mars more
in detail than in earlier periods. Aims of this work
were been to realize a geomorphological map of
Ares Vallis and its valley arms (using data acquired
by recent Martian missions) and to attempt to draft
the geological evolution of the valley.
Methodology: 21 HRSC orbit data-sets andmore than 300 MOC Narrow Angle and THEMIS
VIS images are utilized. We observed and mapped
in detail several erosive and depositional features
distinctive of catastrophic floods, such as
streamlined uplands, erosive terraces, giant bars,
pendant bars, and an impressive cataract-like
feature. Superimposed on the catastrophic floods
morphologies, glacial and periglacial features occur:
they consist of ice-contact structures, thermokarstic
depressions, and patterned grounds. Geological
properties of mapped Units and characteristics of
geomorphological processes responsible for their
shaping are proposed. Impact craters counting was
utilized to constrain the geological evolution of the
area.
Results: Investigations outline that large part of
features characterizing Ares Vallis were shaped by
several catastrophic floods emanated from chaotic
terrains. Geomorphological evidence suggests that
catastrophic floods was, at place, more than 500
meters deep; furthermore, they was ice covered,
confirming that the climatic conditions of the planet
were similar to those of present day [10]. The
amount of time intervening among different floods
varies from hundreds to thousand of years. At the
end of each catastrophic flood, ice masses some
hundreds of meters thick grounded on the valley
floor, forming a stagnant dead-ice body.
Catastrophic floods events were followed by
relatively brief periods of warmer-wetter climatic
conditions, possibly triggered by greenhouse effect
of water vapor and carbon dioxide released in the
atmosphere during the catastrophic flood processes.
During these periods, water in equilibrium with the
atmosphere etched thermokarstic depressions andchannels on areas previously sculpted by
catastrophic floods. Water flowing on ice-walled
streams emplaced ice-contact deposits. Finally, ice
wasted mainly by sublimation processes, indicating
that the Martian atmosphere became again to dry-
cold climatic conditions, similar to those of present
day.
References: [1] Baker V. R., and Milton D. J., (1974), Icarus, 23, 27-41. [2] Baker, V. R., (1992), Mars.
University of Arizona Press, Tucson, 483-522. [3]Komatsu, G., and Baker, V. R., (1997), J.G.R, 102, 4151-
4160. [4] Marchenko, A. G., et al., (1998), Solar System Res., 32(6),425-452. [5] Lucchitta, B. K., et al., (1981),
Nature, 290, 759-763. [6] Lucchitta, B. K., (1982), J.G.R., 87, 9951-9973. [7] Lucchitta, B. K., (2001), G. R.
L., 28, 403-406. [8] Costard, F. M., and Kargel, J. S.,(1995), Icarus, 114, 93-112. [9] Costard, F. M., and Baker,
V. R., (2001), Geomorphology, 37, 289-301. [10] Baker,V. R., et al., (1991), Nature, 352, 589-594.
Figure 1. Geomorphological map of Ares Vallis.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ASSESSING MARTIAN ATMOSPHERIC PREDICTABILITY USING A GENERAL CIRCULATION
MODEL AND A RE-ANALYSIS OF MGS/TES OBSERVATIONS, P. Rogberg1, P.L. Read
1, S.R.Lewis
2,
L. Montabone1&2
.1)
Dept of Physics, University of Oxford, OX13PU Oxford, UK.2)
Dept of Physics and
Astronomy, CEPSAR, Open University, MK76AA Milton Keynes, UK. [email protected]
We assess our ability to forecast the time evolutionof the atmosphere on Mars using a general
circulation model (GCM) with an ensemble
forecasting approach. Model forecasts are compared
with a re-analysis by assimilation of atmospheric
temperature and dust measurements, in this case
using data from the Thermal Emission Spectrometer
on board Mars Global Surveyor.
In a previous study, Newman et. al. (2004)
conducted a 'perfect model ensemble' investigation
of atmospheric predictability using a Mars GCM.
They concluded that the period from northern
hemisphere late autumn to early spring in a typicalMars year had the fastest-growing perturbations and
thus was the least predictable. At other times of year
negative growth rates were found, indicating that
perturbations decayed. In the present work we have
pursued this question further, using a GCM
constrained by data assimilation over nearly three
Mars years (from MY24 to MY27, using the
arbitrary numbering scheme of Clancy et. al.). We
find a rapid development of forecast errors in a free
running GCM in the form of a 'climate drift' with a
global error growth rate varying between 0.5 and 1.5
sols, with maxima around northern hemisphere
solstice and minima around winter solstice andspring equinox. The growth rate is close to the
radiative relaxation time scale. We consider the
growth of perturbations in three equal sized latitude
bands, where we find a significantly stronger error
growth at latitudes poleward of 30 degrees,
compared to the equatorial region. Furthermore,
there is significant interannual variability between
the three years studied so far. In order to identify the
underlying reason for the divergence between model
prediction and the assimilated re-analysis, we
examine the link with transient weather phenomena,
and discuss the role of deficiencies in the model.
Atmospheric forecasts will become of increasingimportance for Mars in relation to planning future
missions, such as ExoMars, and the use of present
and forthcoming spacecraft data (including Mars
Express and proposed missions such as the Mars
Environment and Magnetic Orbiter (MEMO)) in
testing and constraining atmospheric general
circulation models is vital. The use of such models,
in combination with data assimilation, to establish
entry, descending and landing (EDL) parameters for
future entry probes and landers may become
operational once the atmospheric predictability has
been assessed for different seasons and locations.
This would provide a guide to the likely reliabilityof a forecast made in advance of an EDL procedure,
if sufficient data coverage is provided by an orbiting
satellite
Figure 1 shows the rate of the global error growth,, between the assimilated data and control run
during nearly three Mars years. The growth rate is
estimated by a linear fit during the first sol, the
shaded area shows the standard deviation around the
fit.
References:
Newman, C.E., Read, P.L., Lewis, S.R. (2004)
Q.J.R. Meteorol. Soc. 130 DOI 10.1256/ qj.03.209.
Clancy, R.T., Sander, B.J., Wolff, M.J., et.al. (2000) J.
Geophys. Res. 105 pp 9553-9571
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ASSOCIATIONS BETWEEN WATER AND MINERALS ON THE MARTIAN SURFACE AS SEEN
BY VISIBLE AND NEAR-INFRARED SPECTROSCOPY. A. Pommerol1, B. Schmitt
1, J.-P Bibring
2and
the OMEGA Team.1Laboratoire de Planétologie de Grenoble, BP 53, 38041 Grenoble Cedex 9 France.
2IAS,
Université Paris 11, Orsay, France. [email protected]
Introduction: Visible and Near-infraredspectroscopy is used to characterize the composition
of the Martian surface from Earth and from the
Martian orbit. The spectral range of near infrared (1-
5 μm) is particularly useful to study different types
of associations between water and minerals thanks
to strong absorption bands due to different H2O
vibration modes. Two imaging spectrometers
currently in orbit around Mars (OMEGA/Mex and
CRISM/MRO) are used to map the strength and
shape of these bands, respectively on a global scale
and with high spatial resolution. Global maps of
mineral hydration have been produced using both
the 1.9 μm and 3 μm absorption bands strength
measured by the OMEGA instrument ([1],[2],[3]).
Various properties of water ice in polar regions can
also be investigated thanks to the OMEGA dataset
([4],[5]). However, because bidirectional reflectance
spectra are influenced by a number of different
parameters (composition, texture, measurement
geometry), it is often challenging to extract
quantitative information from reflectance data.
Therefore, numerical modeling of radiative transfer
and laboratory experiments on planetary analogs are
sometimes crucial to interpret remote sensing
datasets.
Methods and Results: We have designed and
built an original experimental facility to measure
bidirectional reflectance spectra of Martian surface
analogs (minerals and water ice, pure or associated)
under conditions representative of the Martian
surface (temperature, water vapor partial pressure).
Our experimental data can then be directly
compared with OMEGA and CRISM datasets.
When this is possible, we compare experimental
results with results from radiative transfer
simulations ([6]). We have studied the effects of
albedo, particle size and measurements geometry
variations on the water-of-hydration bands strength
([7]). We then applied these results to the OMEGA
dataset in an attempt to identify the origins of the
spatial variations of these bands (see discussions by
[2] and [3]). Study of the surface texture effects on
band strength requires the use of complementary
data such as thermal inertia maps obtained by
TES/MGS ([8]). As an example, figure 1 presents
values of the integrated 3 μm hydration band plotted
versus reflectance in the spectrum continuum (from
OMEGA) and thermal inertia (from TES) for one
particular OMEGA orbit. These kinds of diagrams
highlight the complex relationships between band
strength, surface texture and surface albedo. Effectsof measurement geometry variations (incidence,
emergence and phase angles) are studied in a similar
way. We will discuss these results and propose
solutions to identify the origins of the observed
spatial variations and to retrieve the real variations
of the materials water content.
Figure 1. Diagram showing values of the integrated 3 μm
water-of-hydration absorption band versus reflectance in
the spectrum continuum and thermal inertia for one
particular OMEGA orbit.
Quantitative effects of such variations of
materials water content on reflectance spectra are
experimentally investigated thanks to controlled
adsorption of water onto analog materials in
conditions representative of the Martian surface.
The same facility is used to study various types of
associations between minerals and water ice
(deposition of frost on dust, sublimation of a water
ice saturated soil…). We are then building a spectral
database of minerals and water associations that we
use to study the evolution of the Martian polar
terrains in a seasonal timescale. We will especially
detail some analysis of the evolution of water (frost
and hydration) at northern latitudes during early
spring.
References: [1] Bibring, J.-P. et al. (2006), Science 312,
400-404. [2] Jouglet, D. et al. (2007), JGR 112, DOI:
10.1029/2006JE002846. [3] Milliken, R.E. et al. (2007),
JGR 112, DOI:10.1029/2006JE002853. [4] Douté S. et al.
(2007), PSS 55, 113-133. [5] Schmitt, B. et al. (2006) 4th
International Mars Polar Science and Exploration conf.,
Abstract #8050. [6] Douté, S. and Schmitt, B. (1998), JGR
103, 31367-31390. [7] Pommerol, A. et al. (2007) LPSC
XXXVIII, Abstract #1774. [8] Putzig, N.E. et al. (2005),
Icarus 175, 325-341.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ATHMOSPHERIC ESCAPE FROM MARS DURING EVOLUTIOARY TIME SCALES H. Lammer1,
N. Terada2, Yu N. Kulikov
3, H. I. M. Lichtengger
1,1Space Research Institute, Austrian Academy of Sciences,
Schmiedlstr. 6, A-8042 Graz, Austria,2National Institute of Information and Communications Technology,
Tokyo, Japan, Polar Geophysical Institute,3Russian Academy of Sciences, Khalturina Str. 15, 183010
Murmansk, Russian Federation
The evolution of the Martian atmosphere with
regard to its CO2 and H2O inventory is expected to
be strongly affected by thermal and non-thermal
atmospheric loss processes of the lightest neutral
and ionized constituents into space as well as by
chemical weathering of the planetary surface
material. The escape processes depend on the
intensity of the solar X-ray and EUV (XUV)
radiation and on the solar wind density during the
Solar system history. In order to investigate the
evolution of the of the CO2-rich Martian
atmosphere, a diffusive-gravitational equilibriumand thermal balance model is used which allows to
study the heating of the thermosphere by
photodissociation and ionization processes due to
exothermic chemical reactions and cooling by CO2
IR emission in the 15 μm band for different solar
radiation exposures. For reconstructing the Sun's
radiation and particle fluxes from present time to
4.6 Gyr ago, data from the observation of solar
proxies with different ages have been employed.
Based on global 3-D magnetohydrodynamic
(MHD) and test particle simulation models of the
solar wind interaction with the upper atmosphere of
Mars, the loss rate of ions over the planet's history
is estimated. It is further shown how high XUV
radiation fluxes result in a hot and expanded
thermosphere, indicating that the high temperature
of the early Martian thermosphere could have led to
blow-off conditions for neutral hydrogen atoms
even for high CO2 atmospheric mixing ratios.
Finally, the impact of an early planetary dynamo on
the erosion of the Martian atmosphere by the solar
wind is briefly addressed.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THE ATLAS OF MORPHOLOGIC FEATURES OF MARTIAN CRATERS J.F. Rodionova2, A.E.
Jakupova1, G.G. Michel
2, E.N. Lazarev
2.
1Moscow State University, 119991, Leninskie gory, GSP-1, Moscow,
Russia.2Sternberg State Astronomical Institute, 119992, Universitetskij prospekt, 13, Moscow, Russia.
Section 2: The analysis of the data bank of theMorphological Catalogue of Craters of Mars (1) is
fufilled and represented in the maps of the density
distribution of craters of different diameters and
morphology. There are the following maps in the
atlas: The density distribution of the craters in
diameter 10 km and more; The density distribution
of the craters in diameter 10 - 20 km, 20 -40 km, 40
- 80 km, 80 - 160 km, more than 160 km; The
density distribution of craters of the first class of
degradation, second class, third class, fourth class
and fifth class, 12 - 19: The density distribution of
craters with ejecta, terraces and faults, peaks and
circular ridges, hills and ridges, central pits on thecrater floor.
The maps are compiled in the equal area
Molveide projection in a scale 1:80M. The maps are
constructed by taking the number of craters falling
into each 5°5° cell of the planet's surface and
normalized this value to the area of the cell to obtain
a density. We have used the code ArcView 3.1
(ESRI USA) and module Spatial Analist. With the
use of this module the images of the density
distribution of craters have been constructed.
The maps allow to discover an interesting
features of density distribution of craters. For
example the map for craters of the first class shows
that the areas with maximum crater concentration
are disposed along two belts at -25° and +25 ° but
Tharsis Montes interrupt the belt in the northern
hemisphere. It is interesting that there are ten times
fewer craters of the first class to the north and south
of parallel -35° and +35° equally for the younger
plains and volcanic regions in the northern
hemisphere and for the older highland region in the
southern hemisphere. At the same time maps for the
craters of the second and third class show that the
craters of these classes are distributed over all thesouthern hemisphere excluding the area of Argire
Planitia, Hellas Planitia, Hesperia Planum and the
part of The Promethei Terra. The maximum crater
density in the northern hemisphere is in Terra
Arabia and Xanthe Terra with a sharp boundary in
the area of Cydonia Mensae, Deuteronilus Mensae,
Protonilus Mensae and Nilosyrtis Mensae. Most of
the class 4 craters appear to have been obliterated in
the region of Hesperia Planum, Promethei Terra and
Hellas Planitia. Map for the class 5 craters show the
maximum of concentration of these old craters is in
the belt from -55° to -75° and along the zero
meridian.Craters with central pits are placed in areas of
low albedo. The density of craters with fluidized
ejecta have maximum in a band around 25°N in the
regions Chryse, Isidis, Elysium Planitia and the
Solis and Hesperia Planum. Kuzmin (3) noted that
the presence of craters with fluidized ejecta is an
indicator of sub-surface ice-containing rocks, so the
map of craters with fluidized ejecta shows regions
with the highest content of ice in the sub-surface
rocks. The example of the map of the distribution of
the ctaters with ejecta is represented on fig.1. The
maximum of density about 90 craters on the area of
1 million square km is in the Lunar Planum, Chryse
Planitia, Hesperia Planum and in the area of 0 °
meridian to the north of equator.References: [1] Rodionova, J.F. et.al. (2000),
Morphological Catalogue of the craters of Mars. The
Netherlands, 158 p. [2] Michel, G.G., Rodionova J.F.
(2000) Non random distribution of the pits craters.
Abstracts Vernadsky-Brown Microsimposium. [3]
Kuzmin, R.O. (1983) Kriolitosfera Marsa. Moscow, 143
p.
Figure 1. The map of density distribution of craters with ejecta
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Biomarkers for the LD/GC-MS of ExoMars H. Steininger. F. Goesmann. Max Planck Institute for Solar
System Research, Max Planck Strasse 2, 37191 Katlenburg-Lindau, Germany. [email protected]
The composition and the detection of the remains
of extant or extinct life on Mars is an issue since the
Viking missions. As the analysis of pyrolysed soilsamples by a GC-MS on Viking, which was highly
sensitive, showed no sign of organic compounds the
measurement strategy has to change. [1] For this a
comparison to the search for molecules indicating
life here on earth is useful.
The Biomarkers found in sedimentary rock and
oil deposits clearly indicate a biological origin of
organic molecules with a complexity high enough to
exclude non biological formation. These molecules
have to be persistent enough to survive harsh
conditions in the bedrock.
Two prominent groups of biomarkers are the
hopanes and the porphyrins. [2] Both are productsof advanced biochemical reactions indicating that
these molecules occurred late in the development of
biochemical processes. Although the oldest deposits
are several 100 Ma old, hopanes and porphyrins are
still widely used by all forms of life indicating that
no major change in biochemistry happened since
these compounds where deposited.
The molecules found in the sediments have
undergone chemical reactions, for example the
hydrogenation of double bonds, indicating reductive
conditions. It has not yet been proven that the
reducing environment necessary for the formation of the hopanes and porphyrins here on earth has a
counterpart on mars. The proposed aromatic
carboxylic acids as degradation product of organic
material will play a major role if the oxidative
conditions are also present in the bedrock. [3]
For the case of Mars the search for both, relative
large inert molecules and oxidized carboxylic acids,
in the bedrock seems a feasible way to yield organic
molecules as the exact Martian geochemistry is
unknown. The organic molecules are shielded from
decomposing UV radiation because they were
incorporated into the sediment during the formation
of the rocks. The combination of Laser desorptionand pyrolysis will give access to a wide mass range
of organic molecule starting from light pyrolysis
fragments up to larger less volatile organics.
References: [1] Biemann, K. (2007), PNAS 104 ,
DOI:10.1073/pnas.0703732104. [2] Kohnen, M. E. L. et
al. (2000), Geochim. Cosmochim. Acta 59,
DOI:10.1016/0016-7037(95)00338-X. [3] Benner, S. A.
et al. (2000), PNAS 97 , DOI:10.1073/pnas.040539497.
H
H
H H
N
N
N
N Ni2+
Figure 1. Hop-22(29)-ene and nickel-(II)-etioporphyrin.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
BIOSIGNATURES OF ANTICIPATED LIFE ON MARS AND THEIR DETECTION BY MSL AND
EXOMARS Joop M. Houtkooper1, Dirk Schulze-Makuch
2.
1Center for Psychobiology and Behavioral Medicine,
Justus-Liebig-University of Giessen, De Maalkom 7, 1191LP Ouderkerk Amstel, The Netherlands.2School of
Earth and Environmental Sciences, Washington State University, USA. [email protected]
Anticipated Life on Mars: The hydrogen peroxide-water hypothesis [1] offers an explanation for the
hitherto puzzling aspects of the Viking Lander
biology experiments. The lack of detected organics in
the GCMS and the evolution of oxygen in the GEx
experiment point to the possibility of organisms using
a mixture of hydrogen peroxide and water in their
intracellular fluid.
H2O2-H2O mixtures: Such mixtures have several
properties attractive in Martian ambient conditions:
The freezing point of the eutectic is -56.5oC, with the
tendency to supercool below that temperature, H2O2
is a source of oxygen and energy, and the mixture is
hygroscopic, enabling organisms to attract watervapor from the atmosphere well below saturated
conditions.
The compatibility of H2O2 with biochemistry may
be questioned, but even terran biochemistry has a
number of uses for H2O2, as a messenger molecule, as
a defence by antibodies, in the metabolism of
Acetobacter peroxidans, with the most extreme
example known being the 25% H2O2 solution
produced by the insect Brachinus crepitans .
Properties of the anticipated organisms: First, the
organisms are well able to grow without liquid water
at Martian ambient temperatures. Furthermore they
may well have an excess oxidative content and theirability to scavenge water vapor from the atmosphere
may also be a disadvantage: Under water vapor rich
conditions, such as saturation at above zero
temperatures, they may have no defence against
hyperhydration. Obviously, the organisms have to
produce their H2O2 from the atmosphere, of which
the H2O2 content is in the ppb range. Moreover, the
contact between the H2O2 and the proteins in the cell
may require an active stabilization mechanism,
possibly in conjuction with the H2O2 production. The
active stabilization mechanism may be similar in
some way to the damage repair mechanism in Deino-
coccus radiodurans, which requires an active
metabolic state.
Explaining the Viking results: The lack of
detected organics can be explained by the gradual
heating to pyrolysis, by which the organisms
decomposed into CO2, H2O, O2, N2 and little else. In
the GEx experiment, the added moisture caused the
organisms to hyperhydrate and perish as well. Theexcess oxidative content of organisms explains the
evolved O2. Moreover, the presence of organisms
may explain the diminished reponses in the "cold
sterilization" tests and the lack of response of samples
stored for 3 months in the dark at above zero
temperatures, possibly because of energy use by
active cellular stabilization.
The MSL and ExoMars instruments: MSL and
ExoMars may reveal indicators of biology. The MSL
will have the SAM searching for carbon compounds,
both in soil samples and in the atmosphere. The soil
samples will be heated to get the GCMS to work. A
laser spectrometer will analyze isotopic abundances.Environmental monitoring will measure humidity.
ExoMars will carry GEP to monitor the environment
for a few Martian years. The Pasteur package on the
rover will contain a GCMS and a microscopic
spectrometer. Organics and oxidants should be
detected and an antibody-based life-marker chip may
detect present life if a biochemistry similar to ours is
involved.
Adding water is no option: Most biology
experiments use liquid water. The resulting com-
bination of humidity and temperature is unmartian
and may cause the anticipated organisms to perish.
The same may well have happened when heating thesamples in the Viking GCMS. This limits the pos-
sibilities of detection. Still, imaging and microscopic
spectroscopy, in the visual and UV to detect
absorbing/photosynthetic pigments, and in the IR to
possibly detect a Raman signature of H2O2 are
possibilities. However, metabolism should not be
neglected: Organisms which produce H2O2 from H2O
and CO2 using photosynthesis or thermal gradients
have to produce reducing species such as CO, CH4
and CH2O. Scavenging O2 from the little present in
the atmosphere is another possibility. These
metabolic processes may result in signatures in the
surface boundary layer of the atmosphere in the form
of diurnal rhythms. Monitoring of the atmospheric
(isotopic) composition with high precision would be
called for.References: [1] Houtkooper, J. M., and D. Schulze-
Makuch, (2007) IJA 6: 147-152.
doi:10.1017/S1473550407003746.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
CHARACTERIZATION OF RADAR-TRANSPARENT DEPOSITS IN THE SOUTHERN AND
EASTERN ELYSIUM REGION OF MARS A. Safaeinili1, R. Orosei
2. R. Phillips
3J. Plaut
1, K. Doubleday
1,
Y. Gim1, B. Campbell
4, R. Seu
5,1Jet Propulsion Laboratory, Pasadena, 91109, USA.
2Istituto di Astrofisica
Spaziale e Fisica Cosmica, Istituto Nazionale di Astrofisica, Rome 00133, Italy,3Department of Earth and
Planetary Sciences, Washington University,St. Louis, MO 63130, USA,4
Center for Earth and Planetary Studies
Smithsonian Institution, Washington, DC 20013-7012, USA, 5Dipartimento INFOCOM, University of Rome``La Sapienza'', Rome 00184, Italy. [email protected] ,
Introduction: We present results on the Elysium
Planitia near the equator that show evidence of
shallow radio-transparent deposits. This is the
largest radar transparent area away from the Mars
polar layered deposits observed by either MARSIS
(Mars Advanced Radar for Subsurface and
Ionospheric Sounding) and SHARAD (SHAllow
RADar). The first indications of radar transparency
in this region was provided by MARSIS. However,
most of this region remained unexplored partly
because of a lack of depth resolution of MARSISthat is ~80 meters in the subsurface. SHARAD is
able to resolve these deposits better due to its higher
bandwidth providing approximately 10 times better
depth resolution. The MARSIS radar sounder
operates over 4 bands between 1.3 MHz and 5.5
MHz and has a maximum bandwidth of 1 MHz.
SHARAD operates between 15 MHz and 25 MHz
with a maximum bandwidth of 10 MHz. Although
SHARAD has higher resolution, its performance
can be degraded due to surface roughness causing a
loss of coherence in radar echo which in turn can
make the detection more difficult.
Radar Observations: We measure the depth of the
deposit to be up to 200-meters thick assuming a
dielectric constant of 4. The origin of this terrain
has been attributed to volcanic flows [1,5], but also
aqueous or sedimentary processes [2]. If there are
remnant water ice deposits in this region, they must
be covered by a protective layer of material to
prevent sublimation. The question is whether the
radar data can provide clues about the nature of
these deposits. The radar data indicate transparent
deposits that can be consistent with both an aqueous
and volcanic origin. However, these radar data can
also measure the extent of these regions, which can
provide additional information about the nature of
these deposits. The fact that we observe a similar
thickness and distribution of deposit in the Elysium
[3] as well as the Amazonis [4], suggests the
possibility of a common mechanism responsible for
these deposits. We have collected data over many
parallel tracks enabling us to develop a map of the
subsurface of this region.
Figure 1 shows radargrams from three parallel
SHARAD tracks. The radargram shows a
Figure 1. Three parallel SHARAD tracks over the
Elysium planitia showing a single subsurface interface at
a depth of 60 m. continuous interface with clear indication of a
second deeper interface in the beginning of the
track. We interpret these as boundaries between
different flow episodes. It is possible that other
boundaries exist but it is not observed by the radar.
Summary: Both MARSIS and SHARAD provide
evidence of an extensive shallow (< 300 meters)
radio-transparent deposit covering the northern
plains of Mars including the Elysium [3] and
Amazonis [4]. These deposits are up to 200-meters
thick in the Elysium region. Similar depth and radar
signature between the Elysium and Amazonis
regions point to common mechanism.Acknowledgments: SHARAD was provided by the
Italian Space Agency (ASI) for use on NASA’s
Mars Reconnaissance Orbiter. MARSIS is a joint
project of ASI and NASA. Some of the work
described herein was performed at the Jet
Propulsion Laboratory under contract with NASA.
References: [1] Hartmann, W. K. and D. C. Berman, J.
Geophys Res., 105, 15,011, 2000. [2] Murray et al Nature,
Vol. 34, pp 352-355, 2005. [3] Safaeinili et al., Seventh
international Mars Conference, 2007. [4] Campbell et al.,
Seventh international Mars Conference, 2007. [5] Plescia,
J.B., Icarus 164 (2003) 79–95.
Seventh International Conference on Mars 3206.pdf
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THE COMBINED RAMAN SPECTROMETER/ LASER-INDUCED BREAKDOWN
SPECTROMETER ELEGANT BREADBOARD B. Ahlers1, G. Bazalgette Courrèges-Lacoste
1. E.
Boslooper1,
1TNO Science and Industry, Space and Science, Stieltjesweg 1, 2600 AD Delft, The Netherlands.
Abstract: Amongst the different instruments thathave been pre-selected to be accommodated on-
board the Pasteur payload of the ExoMars rover is a
spectrometer for combined Raman and Laser
Induced Breakdown Spectroscopy (LIBS). It is
regarded as a fundamental, next-generation
instrument for organic, mineralogical and elemental
characterization of Martian soil and rock samples.
Raman spectroscopy and LIBS will be integrated
into a single instrument sharing many hardware
commonalities [1]. For science objectives, the
synergy is evident: the Raman spectrometer is
dedicated to molecular analysis of organics and
minerals; the LIBS provides information on thesample’s elemental composition.
An international team under the lead of TNO has
been gathered to design, build and test an Elegant
Bread-Board (EBB) of the combined Raman/ LIBS
instrument. Low mass, size and resource usage were
the main drivers of the instrument’s design concept.
Heart of the instrument is a specifically designed,
extremely compact, spectrometer with high
resolution over a large wavelength range, suitable
for both Raman spectroscopy and LIBSmeasurements.
Apart from the previously mentioned
spectrometer, the breadboard includes lasers,
illumination and imaging optics as well as fibre
optics for light transfer. Measurements will be made
in two different contexts: outside the rover
laboratory (using the rover’s robotic arm) and inside
the rover laboratory. Optionally a microscope/
close-up imager could be integrated in the design
concept.
A summary of the functional and environmental
requirements together with a description of the
optical design and its performance are described in[2]. The combined Raman/ LIBS EBB realisation
and first test results are presented.
References: [1] Bazalgette Courrèges-Lacoste, G., B.
Ahlers and F. Rull Pérez, Combined Raman spectrometer/
laser-induced breakdown spectrometer for the next ESA
mission to Mars (2007) in press, Spectrochimica Acta
Part A. [2] Escudero Sanz, I., B. Ahlers and G. Bazalgette
Courrèges-Lacoste, Optical design of a combined Raman-
LIBS spectroscopy instrument for the ESA ExoMars
mission (2007) submitted, Optical Engineering.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
COMPARISON OF THE H2O AND CO2 ADSORPTION PROPERTIES OF PHYLLOSILICATE-FREE
PALAGONITIC DUST AND SMECTITES UNDER MARTIAN ENVIRONMENTAL CONDITIONS
J. Jänchen1, R.V. Morris
2, D.L. Bish
3, U. Hellwig
1.1TFH Wildau (University Applied Sciences), Bahnhofstraße,
15745 Wildau, Germany.2 NASA Johnson Space Center, 2101 NASA Parkway, Houston, Texas 77058, USA.
3Indiana University, 1001 E 10
thSt., Bloomington, Indiana 47405, USA. [email protected]
Introduction: Palagonitic tephra (basaltic tephra
containing hydrated volcanic glass of basaltic
composition) has received widespread attention in
the planetary literature because many are basaltic
spectral and magnetic analogs of bright martian
surface materials [e.g., 1,2]. The <5 μm size fraction
of such tephras is also a size analog for the bright
martian dust. Here we present the H2O and CO2
adsorption properties of palagonites from the island
of Hawaii and compare their behavior with that for
two smectites. The results are important in light of
recent results concerning the distribution of up to
~10% water-equivalent H and the identification of hydrated minerals in near-equatorial regions on
Mars [3, 4].
Experimental: The H2O and CO2 adsorption
properties of two palagonites HWMK919 (<5 μm
size fraction; Mauna Kea Volcano; phyllosilicate-
free) [1.2] and PA 6-7 (Phala District) as well as
those of a Ca montmorillonite (STx-1) and a
nontronite (NG-1) were investigated by thermo-
gravimetry (TG), isotherm measurements, N2 BET
surface area determination, and X-ray diffraction
(XRD) methods. H2O isotherms were measured
gravimetrically from 255-313 K at 10-4
-10 mbar
with a McBain quartz spring balance. CO2 isotherms
were determined volumetrically (196-293 K, 0.1-
1000 mbar) using a Quantachrome Autosorb-1
instrument. TG was performed on a SETARAM
TG-DSC 111 apparatus (heating rate 3 K/min).
Samples were equilibrated at a p/ps (H2O) of 0.3
prior to the experiments.
Results and discussion: Figure 1 and Table 1
show the results of the TG and the BET
measurements. Sample HWMK919 (<5 μm)
accommodated significantly more H2O than all
other samples if the total mass loss of 0.31 g/g of
HWMK919 is assigned to H2O. This is consistent
with the high BET surface area of this sample.
Figure 1. TG profiles of montmorillonite, nontronite, PA
6-7, and HWMK919 (from top to bottom right). Figure 2 compares the H2O isotherms of
HWMK919 with those for NG-1 at 293 K. The
maximum uptake of H2O up to 10 mbar appears to
be the same, although HWMK919 adsorbed more
H2O at significantly lower p(H2O) compared with
NG-1 (and also STx-1, e.g., [5]). Moreover,
HWMK919 retained about 0.1 g/g “extra” H2O after
the degassing procedure prior to the isotherm
measurements and NG-1 retained much less.
Similar results were found for CO2 (not shown),
although the hysteresis obvious in Figure 2 was not
seen with CO2, and the CO2 adsorption was
reversible. The hysteresis observed for HWMK919
appears to be a kinetic effect (Figure 2; filled
triangles shift towards adsorption branch within 30days). However, this is not the case for smectites,
which exhibit marked hysteresis resulting from
structural phase transitions that are obvious in
controlled-humidity X-ray diffraction data.
These results show that fine, phyllosilicate-free
palagonitic dust (HWMK919) can hold significantly
more H2O than smectites under p(H2O) and
temperature conditions approaching the martian
surface. Palagonite is a geologically reasonable
hydrated phase on the surface of Mars, and its
presence may account, at least in part, for the
observations of heterogeneously distributed elevated
concentrations of water-equivalent H on the martiansurface.Table 1 Results of N2 BET specific surface area determination at
77 K and TG measurements (water desorption capacities, a)
Material BET in m2 /g a in g/g
HWMK919, <5 μm 203 0.31
PA6-7 183 0.20
Nontronite, NG-1 68 0.17
Montmorillonite, STx-1 77 0.15
Figure 2. Water isotherms of HWMK919 (<5 μm) and
NG-1 at 293 K after degassing at 383 K in high
vacuum, filled symbols denote desorption.
References: [1] Morris, R.V. et al. (2000) JGR 105, 1757.
[2] Morris, R.V. et al. (2001) JGR 106, 5057. [3]
Feldmann W.C. et al. (2004) JGR
109 E 09006. [4]Poulet, F. et al. (2005) Nature 438, 623. [5] Jänchen, J. et
al. (2006) Icarus 180, 353.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
COMPARISON OF OMEGA AND MCD SURFACE TEMPERATURES. D. Jouglet1, F. Poulet
1, J.
Girard1, A. Spiga
2, F. Forget
2, M. Vincendon
1, Y. Langevin
1, J.P. Bibring
1, B. Gondet
1.
1Institut
d’Astrophysique Spatiale, Université Paris-Sud, Orsay, France.2Laboratoire de Météorologie Dynamique du
CNRS, Université-Paris 6, Paris, France. Contact : [email protected].
Introduction: The imaging spectrometer MarsExpress / OMEGA observes the Martian surface in
the near-infrared up to 5 m [1]. This gives the
opportunity to derive surface temperatures. Since
OMEGA data cover a large part of the Martian
surface, a systematic comparison can be performed
with global numerical simulations. For that purpose
we use the results from the Global Circulation
Model of the Laboratoire de Météorologie
Dynamique, released in the Martian Climate
Database (MCD) [2]. The goal of this work is to 1)
test the accuracy of the OMEGA surface
temperature determination, 2) improve our
knowledge of the mechanisms influencing surfacetemperature, and 3) detect variations in OMEGA
calibration and test the accuracy of new calibration
functions.
Data Processing: The studied variables are the
surface temperature TOMEGA measured by OMEGA,
the surface temperature TMCD calculated by the
MCD and their difference T.
To obtain the OMEGA temperature, reflective
radiance is subtracted to OMEGA spectra thanks to
an a priori knowledge of the 5 m reflectance
value. The resulting thermal part is fitted by the
best black body to retrieve the temperature [3].
The spatial resolution of the MCD is muchlower than that of OMEGA data, therefore a linear
interpolation is performed to get MCD data for
each OMEGA pixel. Since the MCD is also
discretized in 12 months a year and in 12 periods a
day a linear interpolation is also performed to fit
the acquisition time of OMEGA results.
Global study: The Latitude – Solar longitude
(Lat-Ls) map of TOMEGA exhibits seasonal trends
that are very consistent with the TMCD Lat-Ls map.
Results from the 1650 first orbits reveal a mean
absolute difference of ~5K (for a mean surface
temperature of ~240K). This indicates that the
OMEGA calibration function is satisfactory.
We expect lower OMEGA temperatures because
cold dust in the atmosphere cools the flux emitted
from the surface and received by OMEGA.
However the proportion of lower TMCD is large and
cannot be neglected. Moreover Lat–Ls maps for T
have also been plotted and compared to TES
measure of dust opacity [4] but no clear correlation
was observed. This suggests that the effect of dust
at 5 m is low, which is consistent with a low
optical extinction at 5 m compared to that at 1 m
or 9 m [5].
We also notice that T exhibits slight variations
over single orbits, without any obvious dependence
with other parameters (albedo, altitude, presence of
clouds). Fig. 1 reveals that the T mean value (fornominal data) is subject to periodic variations with
time. Their origin has not yet been determined;
they are either due to a physical variation not taken
into account by the MCD or to instrumental
variations not revealed by the OMEGA response to
the calibration lamp.
Results over OMEGA not-nominal data: The
measure of the OMEGA response to a calibration
lamp reveals that the instrument absolute
calibration evolves with time (for the L channel)
[6]. For data at a non nominal calibration level, the
nominal transfer function is not adapted and the
absolute radiances are corrupted [3]. Such defectsare clearly detected in the T (fig.1, red curve)
during the calibration state transitions.
New transfer functions have been derived for
such calibration levels in previous works [6]. The
method is based on the comparison of two
observations of a same area acquired close in time
but at different calibration levels. New surface
temperatures have been derived with these new
transfer functions and the resulting T are more
consistent with the nominal data (fig.1, blue curve).
This validates the accuracy of the new transfer
functions and will enable us to derive new ones for
non nominal orbits greater than 1650.
Figure 1. Value of T (surface temperature fromOMEGA minus that from MCD) versus the orbit
number. Red: using the nominal calibration function for
OMEGA. Blue: using new adapted calibration functions
when the calibration state is not nominal. Black:
evolution of the OMEGA calibration state (the nominal
level is 1500DN for the detector).
References: [1] Bibring, J.P. et al. (2004), ESA sp1240.
[2] Lewis, S. R. et al. (1999), JGR 104, 24,177-24,194.
[3] Jouglet, D. et al. (2007), JGR 112, DOI:
10.1029/2006JE002846. [4] Smith, M.D. (2006), 2nd
workshop on Mars Atmosphere Modelling and
Observation, Granada. [5] Santee, M. L. and Crisp D.
(1995), JGR 100, 5465-5484. [6] Jouglet, D. et al.(2007), 7
thMars Conf. Pasadena, Abs. #3157.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
A COMPARISON OF BULK WATER ICE CLOUDS IN GM3 WITH MEASUREMENTS OF ICE
CLOUDS FROM SPICAM. Frank Daerden1, Nina Mateshvili
1, Ayodeji Akingunola
2John C McConnell
2,
Jacek W. Kaminski21
Belgian Institute for Space Aeronomy, BIRA-IASB, Ringlaan 3, B-1180 Brussels,
Belgium2Centre for Earth and Space Science, York University, Toronto, Canada, M3J 1P3.
The Global Multiscale Mars Model (GM3)
(Moudden and McConnell, 2005), which is based on
the Canadian operational weather forecast model
has an improved simulation of the water cycle on
Mars compared to earlier GM3 work (Moudden and
McConnell, 2007). This includes phase transitions
between water vapour and bulk ice particles, eddy
and molecular diffusion, gravitational sedimentation
and transport between the polar caps, regolith and
atmosphere. One year of nadir UV (200-310 nm)
measurements from SPICAM have been used to
retrieve cloud and dust optical thickness after
allowing for Rayleigh and aerosol scattering andsurface scattering. For the comparison with GM3 we
focus on the tropical cloud belt (Ls 90-150). The
results indicate that GM3 performs quite well for
much of the time as measured against SPICAM.
References:
Moudden, Y. and J. C. McConnell, Three-dimensional on-
line chemical modeling in a Mars general circulation
model, Icarus, 188, 18–34, 2007. Moudden, Y., and J. C.
McConnell, 2005, A new model for multiscale modelling
of the Martian atmosphere, J. Geophys. Res. (Planets),
110, E04001, doi:10.1029/2004JE002354.
Moudden, Y., and J. C. McConnell, 2005, A new model
for multiscale modelling of the Martian atmosphere, J.
Geophys. Res. (Planets), 110, E04001, doi:10.1029/
2004JE002354.
Figure 1. log optical depth of ice from GM3, (b) SPICAM optical depth for similar conditions as GM3
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Constraining the Atmospheric Escape at Mars F. Cipriani1, F. Leblanc
2., J.J. Berthelier
3, O. Witasse
1
1ESTEC, Noordwijk, Netherlands.
2Service d Aéronomie du CNRS, Verrières-le-Buisson, France.
3CETP,
Saint-Maur-des-Fossés, France. [email protected]
Introduction : The quantification of the
atmospheric depletion of the Red Planet since 3.5Gyr is a key question to be addressed in order tounderstand the planet’s history and the possibilityof an early wet environment conducive to theemergence of life. Current evolutionary models of the Martian atmosphere and solar wind parameterssuffer from the lack of constraints on the involvedescaping fluxes and processes. Mars Express datahave both revealed a limited visible occurrence of minerals related to the action of water at its surface(Bibring et al, 2006) and low escape rates of heavyionic species as O+, O2+, and CO2+ (Barabash etal, 2007), stressing the necessity of investigating
various water reservoirs and other atmosphericescape channels (see for instance MEMO : MarsEscape and Magnetic Orbiter, Leblanc et al, 2007).
Modelling the escape of neutrals from Mars :
Current 3D Hybrid and MHD models are used to
estimate the escape of ions (Modolo et al, 2005, Ma
et al, 2004) and are in rather good agreement with
the data acquired during the successive missions to
Mars. In comparison, escape of species in the
neutral form is clearly underconstrained by lack of
appropriate instrumentation flown on any of the
previous mission to Mars. Current estimates of the
escape channels tend to indicate that the escape isdominated by neutrals by at least one order of
magnitude at present Solar Conditions. Such
estimations have for instance been derived from a
Monte Carlo simulation of the Martian Exosphere
which integrates non-thermal processes as
atmospheric sputtering and dissociative
recombination of ionospheric ions and allows a
consistent comparison of the escape rates derived
from those processes (Cipriani et al 2007). We give
here a brief review of the main results obtained
through this model and compare our figures with
other estimates of neutral and ionic escape at Mars.
A Hot Neutral Analyzer to constrain the
atmospheric loss processes :
Hot Neutrals escape fluxes clearly appear as key
parameters to constrain both the interaction of the
solar wind with the Martian upper atmosphere, and
the water loss through the atmospheric escape
channel over the Martian history. We present here
the principle of a mass spectrometer-energy
analyzer which allows deriving both the
composition of the escaping flux (mass range 1 to
44 amu) and, by measuring the energy distribution
of such neutrals, the processes by which this escapeoccurs. This instrument is based on a high
brightness ion source (sensibility better than 5.10-3
A.torr-1
), followed by an energy analyzer which
allows simultaneous acquisition of the energy
spectra of the ions at all energies (in an energy range
1 to 15 eV). Energy analyzed ions are then mass
separated within a TOF mass spectrometer.
Specifications of the instruments have been derived
from the above mentioned numerical model of the
Martian Exosphere (see Figure 1).
References: Barabash S. et al, Science 315, 501 (2007);
Bibring, J.P. et al., Science 312, 400 (2006); Cipriani et al,
JGR112, E7 (2007) ; Leblanc et al, LPS XXXVIII , 1338
(2007), Ma et al , JGR109, A7 (2004), Modolo et al,
Ann. Geophysicae 23, 433 (2005); Shematovich et al,
SSR41,2 (2007).
.
Figure 1. Integrated fluxes of suprathermal oxygen atoms in a 20 FOV, as a function of the instrument line of sight and
altitude. Panel (a) shows fluxes derived in the case of Dissociative Recombination of O2+ ions whereas panel (b) shows
fluxes derived in the case of atmospheric sputtering. Each curve relates to a polar circular orbit of the instrument at 400km
(crosses), 525km (diamonds), 1050km (squares), 2100km (circles).
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
CORE FORMATION WITH IMPLICATIONS TO THE MARTIAN DICHOTOMY: THERMO-
CHEMICAL CONVECTION IN A SPHERICAL SHELL. K. Stemmer1
and D. Breuer1.
1Institute of
Planetary Research, German Aerospace Center (DLR), Rutherfordstr. 2, 12489 Berlin, Germany.
Introduction: A fundamental problem in theevolution of Mars is the timing and the origin of the
crustal dichotomy. The southern highlands and
northern lowlands of Mars differ markedly in
average elevation [1] and crustal thickness [2,3].
Although it is generally accepted that this crustaldichotomy is one of oldest features on Mars, theexact timing of the dichotomy formation is strongly
debated. The origin of the crustal dichotomy has
variously been related to external [4,5] and internal
processes [6], but none of the proposed formation
mechanisms has been fully convincing. We suggest
that the crustal dichotomy is formed (or initiated)
very early: Simultaneously or shortly after coreformation by a degree-one core formation process.
We present a fully spherical model of thermo-
chemical convection with a temperature-,
compositional- and pressure-dependent viscosity [7]
to study the process of core formation and the onset
of mantle convection.
Core Formation: Meteoritic impacts on the
planet cause melting in a near-surface layer. Very
rapid segregation of iron in the melted zone forms a
large metallic layer at the boundary between the
melted and solid silicates. Due to the higher density
of the metallic layer compared to the cold silicate
mantle, an instable layering is existent in which the
metallic layer tends to sink to the planetary interior.
A gravitational instability must occur at long
wavelength to explain the crustal dichotomy. It is
conceivable that the hemispheric asymmetry at the
surface can be explained by a low-degree core
formation process which generates an early crust.
We have accomplished exemplary simulations of
the core formation process. The spatial scales of the
instabilities seem to be strongly dependent on the
assumed rheology as well as on the thickness and
composition of the metallic layer. A thicker layer
and a larger viscosity contrast favor a longerwavelength of instability [8]. The time scale of the
core formation process is as much important for the
heat budget of the planet as the spatial scales of the
core formation process. Assuming that core
formation by negative diapirism is the only core
forming process, the downwelling velocity for Mars
is derived by [9] with 0.26 m/a, while the whole
process happens within 13 Ma. Our simulation
confirms that the core formation process can be
assumed at least as rapid as it is derived by [9].
Onset of Mantle Convection: If a metallic
diapir does cross the cold proto-mantle rather veryrapidly and keeps a large temperature due to the
gravitational release, a large heat flux is provided as
the base of the cold mantle as soon as the core startsforming. On the one hand the time scale of the onset
of convection is important to estimate the heat
removal from the center of a growing planet on the
other hand the spatial scales of the onset of
convection strongly influence the heat budget of the
mantle. Assuming a degree-one mantle convection
pattern, the surface dichotomy can be explained as a
direct consequence of the interior dynamics. Partial
melting and thus volcanism is then presumably
concentrated on the upwelling hemisphere. The core
formation process can be seen as the initial
condition triggering the mantle convection and willstrongly affect the spatial scales of the flow, at least
in the early evolution after mantle convection starts.
Summary: We have investigated the idea that
the Martian dichotomy originates from a degree-one
core formation process generating an early crust,
which is possibly also supported by the following
initial mantle convection dynamics. The fully
spherical simulation of thermo-chemical convection
with a complex rheology considering a Rayleigh-
Taylor setup confirm previous studies, that the core
formation process is very rapid and happens in the
first 50 Ma [10] or even in the first 13 Ma [9]. Of
key importance to generate a low-degree pattern of
the sinking metallic material is the large viscosity
contrast between the metal and the silicate and
especially the relatively low viscosity of the metallic
material. Furthermore a degree-one core formation
process could generate an initially low-degree
mantle convection pattern, which could also support
a crustal dichotomy. From our findings mantle
convection would start immediately after the low-
degree core formation process, which is contrary
discussed by [11].
References: [1] Smith et al. (1999), Science, 284,1495-1503. [2] Zuber et al. (2000), Science, 287 , 1788-
1793. [3] Neumann et al. (2004), J. Geophys. Res., 109 ,
E08002. [4] Wilhelms and Squyres (1984), Nature, 309,
138-140. [5] Frey and Schultz (1988), Geophys.. Res.
Lett., 15, 229-232. [6] Wise et al. (1979), J. Geophys.
Res., 84, 7934-7939. [7] Stemmer, Harder and Hansen
(2006), Phys. Earth Planet. In., 157 , 223-249. [8]
Parmentier, Zhong and Zuber (2002), Earth Planet. Sci.
Lett., 201, 473-480. [9] Kleine et al. (2002), Nature, 418,
952-955. [10] Solomon et al. (2005), Science, 307 , 1214-
1220. [11] Choblet and Sotin (2000), Phys. Earth Planet.
Int., 119, 321-336.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
CUPIDO: Biochemical and Basic Geophysical Field Study of Mars Orange Team, Summer School Alpbach
2007 Collaboration, presented by D. Heinzeller1, R. Higgins
2.
1Institut für Theoretische Physik und Astrophysik,
Leibnizstr. 15, 24118 Kiel, Germany.2Department of Experimental Physics, National University of Ireland
Maynooth, Maynooth, Co. Kildare, Ireland. [email protected]
Context: Of all the celestial bodies in the solarsystem, Mars has held a unique status in human
interest from the beginning. The current search for
past and present life beyond Earth is primarily fo-
cused on the red planet. The proximity of Mars to
the present habitable zone around the sun would
appear to have granted it many key ingredients for
the formation and evolution of life, at least to our
present understanding i.e. internal activity providing
a heat source and driving an internal dynamo, the
existence of an atmosphere and possibly liquid wa-
ter. Past and present missions to Mars have revealed
promising discoveries like atmospheric CO2 con-
centrations similar to those found on Earth2, watervapor in the atmosphere
7and water channels on the
surface1. Key questions on the internal structure and
especially the habitability of the subsurface, which
may be protected from destructive radiation, still
remain unanswered.
Aims: At the Summer School Alpbach 2007 on
Astrobiology, the Orange Team designed a mission
to Mars which will broaden our understanding of the
planet. For the first time, biochemical analysis of the
subsurface to a depth of 3-5m will reveal trace bio-
markers and examine the habitability potential at
different layers. Concurrent seismological experi-
ments will investigate the internal structure to amuch greater depth than those previously done (e.g.
MARSIS4). Spectroscopic measurements observed
within the scope of natural and controlled conditions
will also contribute to our knowledge of the plane-
tary atmosphere and surfaces composition.
Methods: Based on an existing mission concept
(Mars-966), two scientific stations will penetrate at
separate sites to a depth of 3-5m into the Martian
surface. Onboard, the biochemical package
CHEMOVITA will detect organic molecules, geo-
logical particles and sugars along with investigating
isotopic ratios, oxidation states of molecules, the
acidity of the soil and the environmental conditions
of three different depths, down to 5m, below the
surface.
The scientific stations will also contain highly
sensitive seismometers located at their heads and
hence buried deeply in the Martian ground. Addi-
tionally, two small seismic stations (based on the
Deep Space 2 design5) will be placed close to one of
the scientific stations in a triangular formation
which will set up a seismographic network on Mars
for the first time (Figure 1: Proposed landing sites
for individual components). An artificial impact
with known strength and location will be created by
crashing an impactor into the center of this network.
This will provide a standard measurement for seis-
mographic activity which will continue to be moni-tored for one Earth year.
Observation of the ejected plume of this impact
with a visible-infrared spectrometer onboard the
orbiter will reveal the composition, in particular the
concentrations of water and methane, of the ejected
surface material. Conducted from the orbiter, spec-
troscopic measurements of the atmosphere and sur-
face will continue at least as long as the seismic
experiments.
Conclusions: The CUPIDO mission will search
for extinct and extant life on mars. It will address
key questions on the habitability of the subsurface
and the internal composition of the planet in aunique and multifaceted way. We propose the
mission as a successor to ExoMars3
and expect its
scientific return to be invaluable to the success of
any sample return or manned mission.
Figure 1. Proposed landing sites based on current know-
ledge and mission objectives. Yellow triangles: scientific
stations; red stars: seismic stations; blue circle: impactor.
References:1Ambard, A., Mouginis-Mark, P.J. (2007),
Seventh Inter-national Conference on Mars, held July 9-
13, 2007 in Pasadena, California, LPIC 1353, Abs. #3043.2Cottini, V., Formisano, V., Grassi, D., Ignatiev,
N.I. (2006), in: Second workshop on Mars atmos-
phere modelling and observations (Eds. Forget, F.,et al.), held February 27-March 3, 2006 in Granada,
Spain.3Kminek, G., Vago, J.L. (2004), LPSC , held March
15-19, 2004 in League City, Texas, Abs. #1111.4Picardi, G., Biccari, D., Cartacci, M., and 17 co-
authors (2007), MSAIS 11, 15.5Smrekar, S., Catling, D., Lorenz, R., and 8 co-
authors (1999), JGR 104, 27.6Surkov, Y.A., Kremnev, R.S. (1998), P&SS 46 ,
1689-1696.7Titov, D.V., Markiewicz, W.J., Thomas, N., Keller,
H.U., Sablotny, R.M., Tomasko, M.G., Lemmon,
M.T., Smith, P. H. (1999), JGR 104, 9019-9026.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THE DENSITY AND TEMPERATURES OF THE UPPER MARTIAN ATMOSPHERE MEASURED
BY STELLAR OCCULTATIONS WITH MARS EXPRESS SPICAM
F. Forget1, J.L. Bertaux
2, F. Montmessin
2, E. Quemerais
2, F. González-Galindo
1, S. Lebonnois
1, E. Dimarellis
2,
A. Reberac2
, M.A. López-Valverde3
1Laboratoire de Météorologie Dynamique, IPSL, Université Pierre &
Marie Curie, BP99, 4 place Jussieu, 75252 Paris, Cedex 05, France,2Service d’Aéronomie, IPSL, Verrière le
Buisson, France, 3Instituto de Astrofísica de Andalucía, Granada, Spain, [email protected]
The observation of numerous stars rising or
setting through the Martian atmosphere as seen by
the SPICAM UV spectrometer aboard Mars Express
allows to retrieve the atmospheric density and
temperature from 60 km to 130 km [1, 2]. This part
of the atmosphere was previously almost unknown
since very few measurements were available (a few
entry and aerobraking profiles). Moreover, General
Circulation Model simulations had shown that this
part of the atmosphere should present a very active
and interesting dynamic.We present one Martian year of observations
with a total of 616 profiles retrieved at various
latitudes and longitudes. The profiles are analyzed
in details, and compared to the predictions of a
General Circulation Model (GCM) [3, 4].
We studied the seasonal, diurnal and spatial
variations. The atmospheric densities exhibit large
seasonal fluctuations mostly due to variations in the
dust content of the lower atmosphere which controls
the temperature below 50 km, and thus the
atmospheric scale height (Fig.1). In particular, the
year observed by SPICAM was affected by an
unexpected dust loading around Ls=130° whichinduced a sudden increase of density above 60 km.
The diurnal cycle could not be analyzed in details
because most data were obtained at nighttime,
except for a few occultations observed around noon
during northern winter. The corresponding mean
profile slightly differ from the mean profile obtained
at the same locations around midnight, and the
observed differences are consistent with propagating
thermal tides and variations in local heating in the
upper atmosphere (Fig. 2).
Comparison with GCM simulations help to
explain the variations. However, the observed
temperatures are found to be significantly colder
than predicted by the GCM above the 0.01 Pa level
(~90 km altitude). The homopause is higher and
colder than expected (Fig.3). In some locations and
seasons, especially during southern summer,
temperature profiles with homopause temperatures a
few kelvins below the CO2 condensation
temperatures are detected, confirming the possible
presence of CO2 ice clouds in the upper martian
atmosphere at low latitudes [5]. .References: [1] Bertaux et al., JGR 111, CiteID
E10S90 (2006) [2] Quémerais et al. JGR 111, CiteID
E09S04 (2006) [3] Forget et al. JGR. 104 , 24,155-24,176
(1999) [4] Gonzalez-Galindo et al., this issue. [5]Montmessin et al., this issue.
Figure 1. Seasonal cycle of the density at 100 km
observed by SPICAM.
Figure 2. Two mean temperature profiles obtained at
about the same latitude and season, but different local
time, illustrating the diurnal cycle
Figure 3. An average of SPICAM temperature profiles
compared to GCM predictions for various dust and EUV
conditions.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
DEPOSITS OF PHYLLOSILICATES IN TERBY CRATER (HELLAS REGION, MARS) FROM
MULTI-DATASETS (OMEGA/MEX, THEMIS, MOC and HIRISE). V. Ansan1, D. Loizeau
1, N.
Mangold1, Ph. Masson, A. Gendrin
2, S. LeMouelic
3, F. Poulet
2, B. Gondet
2, Y. Langevin
2, J-P. Bibring
2, G.
Neukum4
and the OMEGA co-investigator TEAM,1Lab. IDES-UMR8418, bât. 509, Université Paris-Sud,
91405 Orsay cédex, France,2Lab. IAS-UMR8617, bât. 121, Université Paris-Sud, 91405 Orsay cédex, France,
3Lab. Planétologie et géodynamique, 3 rue de la Houssinière, BP 9220, 44322 Nantes cédex 3, France. 4FU,Berlin, Allemagne. [email protected]
Section 1: Terby impact crater is located at the
northeastern part of Hellas region (75°E – 30°S) on
the cratered highlands. Using multi-dataset available
on this area, we investigated the geological story of
Terby crater. MOLA altimetry [Smith et al., 1999]
shows that Terby displays an anomalous
morphology compared to other impact crater of
~200 km in diameter. Instead of a circular
depression with a central peak, it displays an inner
flat topography locally eroded. The good spatialresolution (few m to 100m/pixel) of MOC [Malin et
al., 1998], HRSC [Neukum et al., 2004; Jaumann et
al., 2007] and THEMIS [Christensen et al., 2003]
images improves the geomorphic analysis of
geological features. In addition, the stereo HRSC
images allow to generate a Digital Elevation Model
(DEM) in the central N-S strip of Terby, using the
photogrammetric software developed both at the
DLR and the Technical University of Berlin
[Scholten et al., 2005], with a spatial resolution of
15m/pixel, and vertical accuracy of 6.1 m. The
spectral data acquired by the imaging spectrometer
OMEGA [Bibring et al., 2005] give information
about the mineralogy of the surficial centimetric
layer.
The northern inner part of Terby crater displays
a 2 km thick series of layers which the THEMIS IR
images (100 m/pixel) show that the flat top consists
of a 100 m thick gray layer covering a series of
bright layers in alternance with dark layers. The
visible THEMIS images (18m/pixel) allow to show
that bright layers are sub-horizontal with a constant
thickness of few meters. At the same scale
(15m/pixel), the HRSC nadir image allows to
observe the central part of Terby without problemrelated to the mosaic of images: Layers show a
progressive variation of dips from rim to reach sub-
horizontal dip in the center of Terby crater. The
mosaic of 87 MOC images (1.5 to 6 m/pixel) shows
the detailed geometry of the bright layers. Locally,
they are disturbed by stratigraphic unconformities
between which bright layers exhibit a ~5° dip
southward. At a greater scale, the HIRISE images
allow to observe the recent degradation of bright
layers, with aeolian erosive flutes and yardangs in
several directions, and fracture networks due to
temperature varitions. In addition, some layers are
covered by black dunes. During the two first years
of European mission, the OMEGA spectrometerobserved Terby crater three times at high resolution,
(~300m/pixel, orbits #232, 2316 and 2327). These
orbits display broad absorption band characteristic
of pyroxene signature. Their spatial distribution
corresponds to the flat floor of depression, some
parts of plateau and localized areas on bright layers
corresponding to black dunes observed in HIRISE
images. Only the orbit #232 displays subtle
absorption bands at 1.9 and 2.3 μm in very localized
areas. The 1.9 μm absorption band indicates that
material would be hydrated and its combination
with the 2.3 μm drop would be consistent with
hydrated mineral, e.g. phyllosilicates [Poulet et al,
2005], which would be in good agreement with the
geomorphic analysis [Ansan et al., 2005].
This suggests that bright layers could correspond
detritic sediments eroded by strong winds, and
locally covered by black dunes of pyroxenes.
References: Ansan, V. et al (2005) LPSC XXXVI,
Abstract#1324. Bibring, J-P. et al. (2005) Science, 307,
1576-1581. Christensen, P. R. et al. (2003) Science, 300,
2056-2061. Jaumann, R et al. (2007) PSS
55,doi:10.1016/j.pss.2006.12.003. Malin, M. C. et al.
(1998) Science, 279, 1681-1685. Neukum, G. et al. (2004)
ESA Special Publication. SP-1240. Poulet, F. et al.,(2005) Nature doi:10.1038. Scholten, F et al. (2005).
Photogram. Eng. Remote Sens. 71 (10), 1143-1152. Smith
et al. (1999) Science, 284, 1495-1503.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
DETECTING CH4 AND OTHER TRACE SPECIES ON MARS WITH A SOIR INSTRUMENT. F.
Montmessin1, A.C. Vandaele
2, E. Neefs
2, J.-L. Bertaux
1, and F. Daerden
2,
1Service d'Aéronomie, réduit de
Verrières, 91371 Verrières le Buisson, France,2Belgian Institute for Space Aeronomy, 3 av Circulaire, 1180
Brussels, Belgium. [email protected]
SOIR (Solar Occultation InfraRed spectrometer)is currently part of the SPICAV/SOIR instrument
[1] on board the Venus Express orbiter. SOIR, an
Echelle spectrometer using an acousto-optic tunable
filter (AOTF) for the order selection, is probing the
atmosphere by the technique of solar occultation,
operating between 2.2 and 4.3 μm, with a resolution
of 0.15 cm-1
. This spectral range is suitable for the
detection of several key components of planetary
atmospheres, in particular that of CH4, using the 3
vibrational band located near 3.3 μm.
Detection of CH4 in the Mars atmosphere has
already been reported [2] using the PFS instrument
on board Mars Express and from ground-based
telescope observations [3]. However, due to the
resolution of PFS instrument (1.4 cm-1
), the P- and
R- branches could not be resolved and the detection
relied on the observation of a sharp feature
attributed to the Q-branch of the band. However,
recent measurements [4,5] have revealed the
presence of 16
C12
O18
C lines in the vicinity of the 3
CH4 band, around 2982 cm-1
, which might impair
the CH4 detection.
With its high resolution capability and its high
signal-to-noise ratio (since observing the sun
directly), the SOIR instrument could resolve
methane individual lines in the 3 vibrational bandregion and provide a very robust CH4 detection. In
addition, it could detect methane isotopomers (such
as13
CH4) and other methane photochemical by-
products (C2H6, H2CO…), thereby providing
important clues on the photochemical cycle of
hydrocarbons on Mars.
We will present simulations of spectra such as
would be recorded by a SOIR instrument probing
the Mars atmosphere, allowing the determination of
lower limits for the detection of trace constituents.
A sensitivity study has also been performed
regarding some of instrumental characteristics of theSOIR instrument (band pass of the AOT filter, order
separation of the Echelle spectrometer).
References: [1] Bertaux, J.-L., et al. (2007), PSS (in
press). [2] Formisano, V. et al. (2004), Science 306 , 1758-
1761. [3] Krasnopolsky et al. (2004), Icarus 172, 537-547
[3] Villanueva, G. et al., Icarus (submitted). [4] Vandaele,
A.C. et al. (2007), 39th
DPS annual Meeting, Abs. #45.03;
Bertaux, J.-L. et al., Icarus (submitted)
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
DETECTION OF MINOR AND TRACE ELEMENTS BY LASER-INDUCED BREAKDOWN
SPECTROSCOPY UNDER SIMULATED MARTIAN CONDITIONS
S. G. Pavlov1, H.-W. Hübers
1, R. Preusker
1, I. Rauschenbach
2and E. K. Jessberger
2
1Institute of Planetary Research, German Aerospace Center, Rutherfordstraße 2, 12489 Berlin, Germany.
2
Institut für Planetologie, Wilhelm-Klemm-Straße 10, 48149 Münster, [email protected]
Introduction: Methodical investigations of
Martian analogue materials under simulated Martian
conditions are carried out with a Laser Induced
Breakdown Spectroscopy (LIBS) spectrometer at
the DLR Institute of Planetary Research. The main
objective is to deliver useful input for the calibration
of the combined Raman and LIBS spectrometer,
which is part of the scientific rover payload of the
ExoMars mission, to be launched in 2013 [1, 2]. A
number of detailed studies concerning different
parameters influencing the LIBS signal detectionunder Martian conditions have been published [see,
for example, 3-6]. Because of the limitations of
mass and energy, the LIBS spectrometer on the
ExoMars mission will have to operate at low
repetition and ablation rates and it is probably not
possible to provide the same spectral resolution as in
laboratory conditions. The presented study focuses
on identification of minor and trace elements in the
LIBS spectra from volcanic rock materials under
Martian conditions and for relatively low laser
excitation energies.
Experimental: The LIBS spectrometer in DLR-
Berlin is capable to carefully reproduce theexperimental conditions expected on the Mars
surface and for the LIBS instrument. The special
chamber keeps a 6-7 mbar of a mixed Martian-like
gas atmosphere (CO2, N2, Ar, O2) and temperature
range of 220-290 K. The setup uses a Q-switched
Nd-YAG laser (Continuum Inlite II-20) at 1064 nm
with a repetition rate of 10 Hz, pulse duration of
8 ns and maximum output power of 250 mJ. The
laser was focused on the sample surface in a spot of
70-80 μm. The Aryelle-Butterfly spectrometer
(LTB-Berlin) utilizes a broad band (171-372 nm &
275-898 nm) high resolution (9400-14000) Echelle
monochromator and a gated ICCD camera (Andor).Results: We analysed a few Martian-like basalt
rock samples and also samples pressed from the
powder of the same material at the Martian
atmosphere, 7 mbar and 213 K. Many minor and
trace elements become to be difficult identified at
laser fluencies below 100 J/cm2
for a single (or a
few) excitation pulses or at below 10 J/cm2
for LIBS
signals integrated over 30-50 laser shots. The signal-
to-noise ratio significantly drops down for the
spectra averaged over less than 10-20 laser shots
(Fig. 1). Relative line intensities for different
elements vary significantly for the rock samples
(Fig. 2). LIBS spectra at Martian conditions
demonstrate a non-negligible background signal
decaying at times longer than a few hundreds of ns.
Figure 1. Examples of the LIBS spectra for the Basalt-
Vogelsberg rock samples for different number of the laser
excitation shots. ICCD delay and gate 300 ns and 50 μs.
References: [1]. See the ESA’s homepage for the
AURORA ExoMars mission:
www.esa.int/specials/Aurora/SEM1NVZKQAD_0.html .
[2] Jessberger, E.K. and the International GENTNER
Team (2004), Geophys. Research Abstracts 6, 03878.
[3] Sallé, B. et al. (2005), Spectrochim. Acta Part B, 60,
805. DOI: 10.1016/j.sab.2005.05.007.
[4] Colao, F. et al. (2004), Planet. Space Sci. 52, 117,
DOI: 10.1016/j.pss.2003.08.012.
[5] Cremers, D.A. et al. (2002), LPSC XXXIII, Abs.
#1330.
[6]. Rauschenbach, I. et al. (2007), LPSC XXXIV, Abs.
#1284.
Figure 2. Comparison of the LIBS line intensities for the
Basalt-Vogelsberg rock and pressed from the powder
(density of 2 g/cm3) samples. Excitation power on the
sample surface was 7 mJ. ICCD camera delay was 300 ns
and gate was 50 μs. Signal is averaged over 30 laser shots.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Displacement-length relationship of normal faults in Acheron Fossae, Mars: new observations with
HRSC. E. Charalambakis1, E. Hauber
1, M. Knapmeyer
1, M. Grott
1, K. Gwinner
1.
1Institute of Planetary
Research, German Aerospace Center (DLR), Rutherfordstr. 2, 12489 Berlin, Germany.
Introduction: For Earth, data sets and models haveshown that for a fault loaded by a constant remotestress, the maximum displacement on the fault islinearly related to its length by ld = [1]. Thescaling and structure is self-similar through time [1].The displacement-length relationship can provideuseful information about the tectonic regime. Weintend to use it to estimate the seismic momentreleased during the formation of Martian faultsystems and to improve the seismicity model [2].Only few data sets have been measured forextraterrestrial faults. One reason is the limitednumber of reliable topographic data sets. We usedhigh-resolution Digital Elevation Models (DEM) [3]derived from HRSC image data taken from MarsExpress orbit 1437. This orbit covers an area in theAcheron Fossae region, a rift-like graben systemnorth of Olympus Mons with a "banana"-shapedtopography [4]. It has a fault trend which runsapproximately WNW-ESE.Method: With an interactive IDL-based softwaretool [5] we measured the fault length and thevertical offset for 34 faults. We evaluated the heightprofile by plotting the fault lengths l vs. theirobserved maximum displacement (d max-model).Additionally, we computed the maximumdisplacement of an elliptical fault scarp where theplane has the same area as in the observed case(elliptical model). The integration over the entirefault length necessary for the computation of thearea reduces the "noise" introduced by localtopographic effects like erosion or cratering. We should also mention that fault planes dipping60° are usually assumed for Mars [e.g., 6] and evenshallower dips have been found for normal faultplanes [7]. This dip angle is used to computedisplacement from vertical offset via sinhd = ,where h is the observed topographic step height, and is the fault dip angle.
Depending on the data quality, especially thelighting conditions in the region, different errors canbe introduced by determining the various values. Anerror of 40% in displacement arises if the true dip
angle is only 30°, i. e. if a shallow dipping thrustfault is mistakenly interpreted as normal fault.Based on our experiences, we estimate that the errormeasuring the length of the fault is smaller than10% and that the measurement error of the offset issmaller than 5%. Furthermore, the horizontalresolution of the HRSC images is between 12.5m/pixel and 25 m/pixel, and 50 m/pixel of the DEMderived from HRSC images because of re-sampling.That means that image resolution does not introducea significant error at fault lengths in kilometerrange.
For the case of Mars it is known that in the growthof fault populations linkage is an essential process
[8]. We obtained the -values from selectedexamples of faults that were connected via a relayramp. The error of ignoring an existing fault linkage
is 20% to 50% if the elliptical fault model is used,and 30% to 50% if only the d max-value is used todetermine . This shows an advantage of the ellipticmodel. The error increases if more faults are linked,because the underestimation of the relevant lengthgets worse the longer the linked system is.
Results: We obtained a value of ld = of
about 2102
for the elliptic model and a value of
approximately2
107.2
for the d max-model.
Figure 1. Displacement-length values for 34 normal faultsassuming dip angles of 60°. Red dots mark -values of the elliptic-model, blue dots mark -values of the d max-model.
The data show a relatively large scatter, but they can
be compared to data from terrestrial faults
(22
105...101~
= ; [9] and references therein).
In a first inspection of the Acheron Fossae 2 regionin the orbit 1437 we could confirm our firstobservations [10].If we consider fault linkage, the -values shifttowards lower -ratios, since linkage means that d remains essentially constant, but l increases
significantly.
We will continue to measure other faults andestimate the released seismic moment.
References: [1] Cowie, P. A. and Scholz, C. H. (1992),
JSG, 14, 1133-1148. [2] Knapmeyer, M. et al. (2006),
JGR, 111, E11006. [3] Neukum, G. et al. (2004), ESA SP-
1240, 17-35. [4] Kronberg, P. et al. (2007), J. Geophys.
Res., 112, E04005, doi:10.1029/2006JE002780. [5]
Hauber, E. et al. (2007), LPSC, XXXVIII, Abs. #1338. [6]
Wilkins, S. J. et al. (2002), GRL, 29, 1884, doi:
10.1029/2002GL015391. [7] Fueten, F. et al. (2007),
LPSC, XXXVIII, Abs. #1388. [8] Schultz, R. A. (2000),
Tectonophysics, 316, 169-193. [9] Schultz, R. A. et al.
(2006), JSG, 28, 2182-2193. [10] Hauber, E. et al. (2007),
7th Mars Conference, Abs. #3110
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
DISPLACEMENT-LENGTH RELATIONSHIPS OF NORMAL FAULTS ON MARS: NEW
OBSERVATIONS WITH MOLA AND HRSC E. Hauber1, E. Charalambakis
1, K. Gwinner
1, M. Grott
1, M.
Knapmeyer1, K.-D. Matz
1, M. Wählisch
1.
1Institute of Planetary Research, DLR, Rutherfordstr. 2, 12489 Berlin,
Germany. [email protected]
Introduction: The geometric properties of planetary fault populations provide useful infor-
mation on fractured rock bodies [e.g., 1]. However,
so far only few data on the relationships between
fault length and displacement have been measured
for extraterrestrial faults [2], partly due to the
limited number of reliable topographic datasets.
Here we use MOLA altimetry data [3] and HRSC
images [4] to obtain one or more displacement
values for a given normal fault. This method allows
us not only to measure the maximum displacement,
but also to analyze the displacement distribution
along the trace of a single fault. We compare our
results to previous measurements on Mars and onEarth, and discuss implications for further studies.
Results: We measured 145 faults of the
southeastern Ophir Planum [5] in the Valles
Marineris region (Fig. 1a) fault array (marked in
yellow in Fig. 1b). For each fault, the fault length
and one or more offset (throw) measurements were
obtained in an interactive software tool.
Figure 1. Study area in Ophir Planum, centered at ~9.6°S
and ~292.5°E. (a) Viking Orbiter image mosaic with
physiographic features labeled. (b) Tectonic sketch map.
The en echelon configuration of the two main fault sets
(see [5]) is indicated by different color shading. We
measured 145 faults in the “yellow” fault set.
Displacement-Length Relationship. Thedistribution of maximum displacements ( Dmax) vs.
fault length ( L) appears to be similar to previousmeasurements [6] from the northeastern branch of
the Tempe Terra rift (Fig. 2). However, a tendency
for slightly shorter fault lengths than those obtained
by [6] can be observed. This might be due to the fact
that we measured separate fault segments, since
fault lengths increase in relation to the maximum
offset if linkage is considered [7]. We expect that
D/L values will shift towards lower D/L ratios if we
consider fault linkage in the next step.
Figure 2. D/L values for 145 normal faults on Ophir
Planum. Red dots mark topographic offset, small black
dots mark displacement on fault plane after correction for
60°-dipping fault planes. The data show a relatively large
scatter, but are comparable to data from terrestrial faults
( = ~1-5 10-2; [2]; see also [8]: Dmax = 0.03 L1.06).
Displacement Distribution along Faults. The
displacement distribution along some of the selected
faults has a more or less symmetrical pattern.
However, in many other cases the distribution is
distinctly asymmetrical, an effect that is also
observed for slip distributions at earthquakes on
Earth [9]. We will analyze relay ramps in detail to
determine if asymmetric distributions are an effect
of fault segmentation.
References: [1] Schultz, R. A. (1999) JSG, 21, 985-993.
[2] Schultz, R. A. et al. (2006) JSG, 28, 2182-2193.
[3] Zuber, M. T. et al. (1992) JGR, 97 , 7781-7797.
[4] Neukum, G. et al. (2004) ESA SP-1240, 17-35.
[5] Schultz, R. A. (1989) JGR, 96, 22,777-22,792.
[6] Wilkins, S. J. et al. (2002) GRL, 29, 1884, doi:
10.1029/2002GL015391. [7] Dawers, N. H. and Anders,
M. H. (1995) JSG, 17, 607-614. [8] Schlische, R. W. et al.(1996) Geology, 24, 683-686. [9] Manighetti, I . et al.
(2005) JGR, 110, B05302, doi: 10.1029/2004JB003174.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
DIVERSITY OF SURFACE ROCKS ENCOUNTERED ALONG THE TRAVERSES OF SPIRIT AND
OPPORTUNITY. E. Tréguier1, C. d’Uston
1, P. Pinet
1, J. Brückner
2, R. Gellert
3, and the Athena Science
Team4.
1OMP, 14 av. E. Belin, 31400 Toulouse, France.
2Max-Planck-Institut f. Chemie, Mainz, Germany.
3Dep. Physics, Univ. of Guelph, Guelph, On, Canada.
4Cornell University, NY, USA. [email protected]
Along their ~10 km traverses, the MarsExploration Rovers encountered a large diversity of
rocks [1,2]. A classification based on
multidimensional analysis of the chemical
compositions, determined by both APXS, has been
produced and appears to well characterize this
diversity. Using this unsupervised analysis several
rock classes could be defined at both landing sites,
consistent with their geographic localization (Fig 1).
During the first part of its traverse at Gusev crater,
rover Spirit encountered only olivine-bearing
basalts, relatively homogeneous in composition. In
contrast, the rocks in the Columbia Hills showed
various alterations providing a surprisingly largevariety of rock classes in a relatively small area.
Home Plate is a light-toned subcircular plateau with
exposure of layered bedrock in the Columbia Hills.
It is thought to be the result of volcanic interactions
with water or ice [3]. Rocks encountered by rover
Opportunity at Meridiani exhibit less diversity, butprovide evidence for a diagenesis that involved
water. With the exception of some exotic samples
and meteorites, rocks are layered sandstones that
could be described by the mixture of two
components: siliciclastic material and sulphates. The
bedrock is remarkably constant along the traverse
with distinct stratigraphy inside impact craters. The
negative correlations of several chemical elements
with S (notably for Si, Al, Na, and K) reveal the
siliciclastic fraction, while the positive correlations
(notably Ca and Mg) prove the presence of
sulphates containing these elements [4].
References: [1] Arvidson, R. et al. (2006) JGR 111,
E02S01, doi:10.1029/2005JE002499. [2] Squyres, S. et al.
(2006) JGR 111, E12S12, doi:10.1029/2006JE002771. [3]
Squyres, S. et al. (2007), Science, 316, 738. [4] Brückner,
J. et al. (2007), 7th Conf. Mars, Abs. #3120.
Figure 1. Maps of Gusev and Meridiani with the localization of rock samples until sol 1170 (all rock samples for Gusev;
only abraded samples for Meridiani). The colour code indicates to which class/cluster the samples belong.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
DUST AND WATER ICE CLOUDS IN MARTIAN ATMOSPHERE FROM THE PFS MEX DATA L.
Zasova1, D. Grassi
2, M. Giuranna
2, V. Formisano
2, N.Ignatiev
1.
1Space Research Institute, RAS, 117997
Moscow, Russia.2INAF, Istituto di Fisica Spazio Interplanetario, 00133 Rome, Italy. [email protected]
IR spectra of Planetary Fourier spectrometer (LWC
7 - 35 m) allow to obtain in self consistent way thetemperature profile and aerosol opacity from single
spectrum and consequently temperature field vs.
latitude-altitude and aerosol opacity along each
orbit. We discuss a behavior of temperature and
aerosol opacity (dust and water ice clouds) in the
areas of Tharsis Volcanoes, Hellas and in Valles
Marineris.
In Valles Marineris the water ice clouds was found
in the morning around 10h with opacity of 0.1 - 0.2
(at 825 cm-1) and typical particle size 2 - 4 m at
Ls=24°. (3 - 13°S, 286°E). A surface temperature
was found about 10K lower than outside of Valles
Marineris and it remained lower up to 20 kmaltitude. In the same area in Valles Marineris the
observation at Ls=135°, LT=13h showed the dust
storm with opacity of 1. It was composed of silicate
dust, typical for Martian dust storms. Dust opacity
was found maximal in Valles Marineris. Another
example is observation at LS = 38 deg (13-15°S,
302E) of the morning haze. The haze was observed
simultaneously by PFS, OMEGA and HRSC. From
PFS LWC the opacity was estimated to be about
0.3. This haze may consist from water ice with
mean particle radius exceeding 4 m or two layered
particles consisting of ice and dust. It is difficult to
explain this haze by another dust composition,because dust on Mars is mixed well and silicate dust
is definitely observed in Valles Marineris and
identified by pronounce band at 1075 cm-1.
Orbit 41 through Hellas
L a t i t u d e , d e g
O p a c i t y a t 1 0 7 5 c m - 1
T a u ( 1 0 7 5 )
Hellas
O1109, Ls=122
O41, Ls=338
Hella
s
Surface temperature
Tau(825 cm -1)
T s u r f ,
K
T s u r f
terminator
Figure 2. Hellas in two seasons.
In Hellas well mixed dust was observed at
Ls=338° near noon, along whole orbit with average
scale height of 11.5 km and opacity in Hellas up to
1. Water ice fog with opacity of several units was
observed in Hellas at Ls=122° in the near surface
layer below temperature inversion in the
atmosphere.
The ice clouds are observed in all seasons above
Tharsis volcanoes with increasing opacity to
several units at northern summer, when equatorial
cloud belt is also observed. Opacity and particle size
increase in the afternoon.
L.Zasova and N. Ignatiev acknowledge the RussianFoundation of Basic Research for financial support,
grant 07-02-00850.
Figure 1. Morning haze in Valles Marineris
Hsurf/4
O438, Ls=38
Tau(825)
Tau(1075)
Equat. Cloud belt O p a c i t
y
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DYNAMIC MARS: NEW DARK SLOPE STREAKS OBSERVED ON ANNUAL AND DECADAL TIMESCALES. K. S. Edgett
1, M. C. Malin
1, R. J. Sullivan
2, P. Thomas
2, J. Veverka
2,
1Malin Space Science Systems,
P.O. Box 910148, San Diego, CA 92191-0148, USA;2CRSR, Space Sciences Building, Cornell University,
Ithaca, NY 14853, USA.
Introduction: Dark slope streaks are long, some-
times tapered features common to steep slopes in thehigher albedo equatorial regions of Mars. They werefirst observed in some of the highest resolution Vikingorbiter images [1Ð3] and numerous subsequent exam-
ples have been found in high resolution (1.4Ð20m/pixel) Mars Global Surveyor (MGS) Mars Orbiter Camera (MOC) images taken since September 1997.Sullivan et al. [4] summarized first-year MOC obser-vations of dark slope streaks and possible formationmechanisms. In general, they appear to result frommass movement, but details regarding the particlesizes, nature of their flow, and initiation mechanismsare still under study. The purpose of this paper is to
present an exciting new observation that will eventu-
ally lead to an understanding of the rate at which dark slope streaks form, and the rate at which they disappear (thought to be via dust mantling). This paper reportstwo cases where the same location was photographedtwice, and new dark slope streaks were observed in themore recent images. The first case shows changes thatoccurred between 1978 and 1999 (~1 Mars decade), thesecond case shows changes that occurred between early1998 and late 1999 (~1 Mars year).
Change in 11 Mars Years: On 7 May 1978 (Ls
~83¡), the Viking 1 orbiter acquired a ~17.5 m/pixelimage of an 11.6 km diameter impact crater locatedwithin the Schiaparelli Basin at 1.8¡S, 343.9¡W. A
portion of this image, 748A12, is shown in Figure 1a.This was one of the various examples from the Vikingmission that showed the existence of dark slope streakson Mars. The crater was next seen on 15 August1999ÑLs ~187¡, about 11.3 Mars years later (MGSMOC image M04-01105 in Figure 1b), and againabout 11.5 Mars years since the 1978 image at L s
~246¡ on 18 November 1999 (M09-04689 in Figure1b). The arrows in Figure 1b indicate the streaks thatwe can confidently identify as being new in 1999 rela-tive to the 1978 image. Some of the streaks seen in1978 appear to remain in 1999, perhaps indicatingeither that the rate at which streaks fade or becomeobscured is slower than the rate that new ones form or that certain slopes are sites of repeated streak formationover the course of 11 martian years. Figure 1 showsonly the northeast quarter of the crater at 1.8¡S,343.9¡W, and the MOC images indicate at least 7 newstreaks occurring on the crater wall and at least 6 onslopes outside the crater rim. To first order, one mightconclude that this is approximately 1 new streak per year for this quarter of the crater, or perhaps 4 newstreaks per year if the entire crater was in view. Regard-
less, the images in Figure 1 indicate a planet upon
which mass movements occur today in the modernmartian environment.
Change in 0.9 Mars Year: The second example inwhich new dark slope streaks are observed is a muchmore dramatic example than that shown in Figure 1.Instead of documenting changes that occurred since theViking missions, Figure 2 shows a case in which newslope streaks formed in less than 1 martian year andwere photographed exclusively by the MGS MOC.Figure 2a is a subframe of a MOC image taken on 1February 1998 at Ls ~266. Figure 2b shows the samelocation on 18 November 1999 at L s ~246, less than 1Mars year later. The location is the southeastern quarter
of an impact crater north of Apollinaris Patera at6.0¡S, 183.8¡W. Three new dark slope streaks formedduring the 0.92 Mars year interval between the two
pictures. During this same interval, the older streaksthat were present in February 1998 remained visible.
Discussion: The 1999 MOC images presented herewere specifically targeted to look for changes in thenumber and relative brightness of dark slope streaksover time. New streaks are observed, and it is observa-tions like these that will eventually lead to a better understanding or determination of slope modificationrates in the modern martian environment. Differencesin spatial resolution and illumination conditions be-tween older and newer images can make quantificationof these changes difficult, but for the moment what isinteresting is the fact that changes can be documentedat all. The presence of new slope streaks is a key indi-cator that geologic processes other than wind action arein fact at work on Mars today. We also note that thenewer streaks (Figs. 1b and 2b) are darker than older streaks seen in earlier images, consistent with the hy-
pothesis [e.g., 1Ð4] that dark slope streaks fade withtime. The appearance of a new dark streak is probably asudden, catastrophic change, while the disappearance of an older streak may be a more gradual process becauseolder streaks remain visible over annual and possiblydecadal time scales.
References: [1]ÊMorris E. (1980) JGR, 87 , 1164-1178. [2]ÊWilliams, S. H. (1991) LPSC XXII , 1509-1510. [3]ÊFerguson and Lucchitta (1984) NASA TM 86246 , 188-190. [4]ÊSullivan R. et al. (1999) LPSC
XXX, #1809.
unar and Planetary Science XXXI 1058.pdf
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NEW DARK SLOPE STREAKS: K. S. Edgett et al.
Figure 1. Appearance of new dark slope streaks over 11 Mars year interval. (A) On the left is a subframe of Vikingorbiter image 748A12 taken in May 1978. White outlines indicate the location of two MOC images obtained in1999. (B) On the right are the two MOC image subframes (M09-04689 from November 1999, and M04-01105 fromAugust 1999) overlain on the earlier Viking image. New dark slope streaks are indicated by arrows. Location is acrater at 1.8¡S, 343.9¡W. North is up, illumination is from the left.
Figure 2. Appearance of 3 new dark slope streaks (arrows) after a 0.92 Mars year interval. Image on left is subframeof MGS MOC AB1-11304, image on right is M09-04872. These are located in a crater at 6.0¡S, 183.8¡W. Bothviews are illuminated from the lower right, north is toward the left.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
EFFECT OF THE WATER AMOUNT INCREASING IN THE UPPER LAYER OF THE MARTIAN
SOIL DURING WINTER SEASON AROUND THE LATITUDES ±50° BASED ON THE TES TI DATA
ANALYSIS. R.O.Kuzmin1, E.V. Zabalueva
1, P.R. Christensen
2
1Vernadsky Institute of Geochemistry and
Analytical Chemistry, Russian Academy of Sciences, 19 Kosygin str., Moscow 119991, Russia, [email protected];2School of Earth and Space Exploration, Arizona State University, Tempe, Arizona, USA;
Introduction: As it has been shown recently, the
essential increasing of the water amount (ice + bound water) in the Martian surface layer was
observed during winter season [1, 2, 3]. In the work
we had conducted the analysis of the TES TI withgoal to investigate the order of the winter-time
increasing of the water ice amount within the
Martian surface layers with thickness in 3-10 cm.
Analysis and results. For the study we
conducted mapping of the TES TI within one of thesectors of Mars (±50°, 230°-300°W) for summer
(Ls=150°-160°) and winter (Ls=300°-310°)
seasons. Both analyzed seasons had lowatmospheric dust opacity [4] that excluded the
influence of the parameters on the thermal inertia
determination. Mapping results (Fig.1a,b)demonstrates that the values of the TI during winter
season in the latitude range 30°-50°N are becoming
notably higher to comparison with summer-timevalues. Apparently, such winter-time increasing of
the TI value has been provoked by appearance of
some water ice (or frost) amount within the surfacelayer. To estimate the possible ice increase in the
soil during the winter we used the nomogram,
created for ice content determination [1] based onrelationship between the TI dry soil and the TI icy soil
values (computed for different soil’s ice content
from 0% to 10%). For this the mapped summer andwinter TI values were zonally averaged (in 5°
latitude belts) and then were plotted on the
nomogram (Fig.1c, d). As it seen from Fig.1c,d, thezonally averaged winter-time TI values are
corresponds to the ice amount of 5-10 vol. % for
latitude ranges 35°-50°N and 40°-50°S respectively.At the lower latitudes (0°-30°N and 0°-35°S) the
winter-time TI values are consistent with much less
soil ice increase (mostly < 2 vol. % and up to dry
soil). To compile the winter-time map of the ice
distribution within surface layer in the studdedsector of Mars we fulfilled next procedure. From all
mapped TES TI data we extracted only those TES
surface footprints (from summer and winter maps),which have geographic location coincidence at
accuracy < 0.05°. The thermal parameters of the
soil-ice mixture for the winter-time footprints wereestimated by similar method as used in [5, 6]. The
two-component mixture is characterized by next
thermal parameters: ρc = ερicecice+ ρdrycdry and k = ε k ice + I 2dry /( ρdrycdry). These two expressions are
substituted into formula of thermal inertia. After
simple manipulation, one can receive the quadratic
equation: aε 2+bε +c=0, where a= ρ icecicek ice,
b= ρ drycdryk ice+ ρ icecicek dry, and c=I 2dry-I 2. At that, I dry
and I represent the thermal inertia values for summer and winter seasons respectively. So, having
the TES TI data for the same place (ϕ ,λ ) from two
seasons we have solved the equations relatively
unknown parameter (ε - ice vol. % ) and compiled
the map of the winter-time ice distribution withinthe studded sector of Mars for both hemispheres
(see Fig.1e). The Fig 1f shows the zonally averaged
(in 5° latitude belt) of ice content as function of the
latitude. The received results have demonstrated
existing of the strong seasonal effect of the water ice (or frost) amount variations in the surface layer
with thickness from ~ 3-10 cm.
Figure1. Regional TES TI maps for summer (a) and
winter (b); t t hhee nomograms for the TI dry soil and TI icy
soil values at different soil ice content with plotted
relationship between zonally averaged (in 5° belts) the
summer- and the winter-time TI values in the Northern
(c) and the Southern (d) winter; e – the map of the
winter-time ice increase in the surface layer of Mars,derived from TES TI data, and it’s zonally averaged
values (f).
References: [1] Kuzmin R.O. et all. (2007), VII MarsConference, Abstract 3022. [2] Kuzmin R.O. et al.,(2005)
LPS XXXVI, Abstract 1634.[5] [3] Kuzmin R.O. et al.
(2007) Solar System Research, 41, 2, 89-102. [4] Smith
M.D. (2004), Icarus, 167 , 148-165. [5] Mellon M.T. and
Jakosky B.M. (1993) JGR, 98, 3345-3364. [6]Schorghofer N. and Arharonson O. (2005) JGR, 110,10.1029/2004JE002350.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
ENERGY RELIEF ANALYSIS OF EASTERN TITHONIUM CHASMA,MARS. D. Baioni1, L. Lanci
1, F.
C. Wezel1.1Institute of Earth Science University of Urbino, Campus Scientifico Sogesta 61029 Urbino (PU),
Italy. [email protected]
A morphostructural investigation in theeasternmost part of Tithonium Chasma trough has
been carried on using the analysis of the energy of
relief. Energy of relief represents the potential
energy defined as result of the product of rock mass,
elevation and gravity acceleration. An analysis of
this parameter is performed by measuring unit cells
of equal areas and assuming that for such cells mass
and gravity acceleration can be treated as constant
values. This quantitative parameter is thus expressed
by the maximum difference in elevation between the
highest and lowest point measured in a given area.
On the Earth energy of relief is used to detect the
intensity of denudation processes on the landscapeand it is used as useful geomorphic marker in the
morphotectonic and morphostructural
investigations. In fact if calculated for single cells of
small size this parameter may reveal zones
characterized by tectonic movements that controls
the development of relief.
The distribution of the energy relief in the study
area was obtained by subdividing the study area into
square cells of 1 km2, which were numbered by
orthographic coordinates. Within each cell, the
value of the energy relief was calculated. Highestand lowest elevation for each cell was obtained from
the topographic map with 200 m contour lines. This
procedure gave a detailed map of 4840 values of
energy of relief parameter in the entire study area.
The values obtained were then used to construct of a
map of energy relief distribution.
The analysis of distribution of the energy of relief
highlights anomalies such as different values
between opposite chasma walls, a general increasing
trend of values toward north and two sets of
lineaments of high energy of relief that occurs along
the bottom of the chasma.
Our investigation suggests that in the study areathe tectonic activity is “younger” in the north side
then in the south one and that the bottom is affected
by lineations, which are interpreted as faults.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
EVIDENCE FOR NON-POLAR ICE DEPOSITS IN THE PAST HISTORY OF MARS. James W. Head1
and David R. Marchant2,
1Dept. of Geological Sciences, Brown Univ., Providence, RI 02912 USA
2Dept. of
Earth Sciences, Boston Univ., Boston MA 02215 USA ([email protected]),
Introduction: The polar caps provide a record of the
recent climate history of Mars [1]. Studies of the spin-axis/orbital parameter history provide a robust solution for
the most recent ~20 Ma of martian history, but cannot be
mapped further back into the past [2]. Thus, deconvolving
the complex climate history of Mars requires analysis of the
basic geological information, and interpretation of the
depositional record of glaciation and glacial conditions at
non-polar latitudes. These interpretations are assisted by
polar analogs to Mars (such as the Antarctic Dry Valleys)
[3]), and an understanding of the behavior of polar ice under
different insolation conditions using GCMs [4-5]. Finally,
the availability of very high resolution images and
topography (e.g., MOLA, MOC, CTX, HRSC, HiRISE)
provides the ability to characterize and interpret these
deposits. We report on recent analyses assessing thepresence, age, and significance of non-polar ice deposits as
evidence of the history of climate on Mars.
(1) Latitude-Dependent Mantle and Recent Ice Ages:
Multiple lines of evidence have been presented that show
the presence of geologically very young and unusual
features and deposits that formed as a result of recent quasi-
periodic climate change [e.g., 6]. Latitude is the single
variable with which all of these diverse observations
correlate, and climate is the only process known to be
latitude-dependent. The very strong correlation between the
nature of the terrain smoothness, the continuity of the
mantle, the high interpreted water content, and the
theoretical stability of ice in the near-surface soil, all
compellingly point to climate-driven water ice and dust
mobility, and emplacement during recent periods of higherobliquity [2]. Degradation and dissection of the deposit in
mid-latitudes point to recent climate change [e.g.,7],
reflecting return of mid-latitude ice to polar regions during
recent lower obliquity [e.g., 6,8]. (2) Northern High
Latitude Cold-Based Glacial Crater Fill: Ridges arrayed
in lobate patterns have been interpreted as drop moraines
deposited during retreat of a lobate cold-based glaciar
originating on the crater rim [9]. (3) Mid-High Latitude
Concentric Crater Fill (CCF): New data show
morphology and structure that support the role of ice in
CCF formation [10], and that CCF craters may have been
ice-filled [11]. (4) Mid-Latitude Lineated Valley Fill
(LVF) and Plateau Glaciation: Earlier studies emphasized
the role of vapor-diffusion-assistedemplacement of ice in slope-related
talus piles, causing talus lubrication
and plastic flow of debris [12]. New
data show that significant ice was
involved and that debris-covered
glacial flow formed regional valley
glacial landsystems [13,14]. (5)
Mid-Latitude Lobate Debris
Aprons (LDA): Earlier thought to
represent ice-assisted creep [12],
LDA internal structure and
morphology now point to debris-
covered glaciers for many [15]. (6)
Evidence for Mid-Latitude Ice
Highstands: New data show
evidence for highstands suggesting that almost a kilometer
of ice has been lost from LVF [16]. (7) Low Mid-LatitudePhantom Lobate Debris Aprons: Former ice-rich deposits
surrounding massifs at latitudes even lower than the LDA
are observed [17]. (8) Tropical Mountain Glaciers
(TMG): New data suggest that fan-shaped deposits on the
NW flanks of Tharsis Montes and Olympus represent huge
TMGs [18] formed during Late Amazonian periods of high
obliquity [19]. (9) Near Equatorial Outflow Channel Rim
Deposits: Glacial-like features on the Mangala rim (18°S)
suggest that climate earlier in the Amazonian was cold in
the near-equatorial regions [20]. (10) South Circumpolar
Ice Cap: The Hesperian Dorsa Argentea Formation
(DAF): Hesperian-aged south circumpolar deposits (DAF)
have been interpreted as a very large volatile-rich polar
deposit; its characteristics (e.g., sinuous ridges interpretedas eskers, marginal fluvial channels, etc.) have been
interpreted to indicate that the DAF contained significant
quantities of water ice, representinng an ancient circumpolar
ice sheet [21].
Summary: Together, these data provide insight into the
climate history of Mars; they suggest that the climate has
been similar to that of today for much of the Amazonian,
with climate variations being driven largely by changes in
spin-axis/orbital parameters [2]. The Hesperian-aged DAF
suggests that conditions were different in this important
transitional period, with the possibility of a thicker
atmosphere, providing an important context for the
assessment of the Noachian climate history of Mars.
References: 1) M. Carr, Water on Mars, 1996; 2) J. Laskar etal, Icarus, 170, 343, 2004; 3) D. Marchant and J. Head, Icarus, in
press, 2007; 4) M. Richardson and J. Wilson, JGR, 107, 5031,
2002; 5) M. Mischna et al., JGR, 108, 5062, 2003; 6) J. Head et al.,
Nature, 426, 797, 2003; 7) J. Mustard et al., Nature, 412, 4211,
2001; 8) J. Laskar et al, Nature, 419, 375, 2004; 9) J. Garvin et al.
MAPS, 41, 1659, 2006; 10) M. Kreslavsky et al., MAPS, 41, 1659,
2006; 11) J. Head et al., Vernadsky-Brown Micro 46 , 2007; 12) S.
Squyres et al., Mars, U of AZ Press, 523, 1992; 13) J. Head et al.
EPSL, 241, 663, 2006; 14) J. Head et al., GRL, 33, L08S03, 2006;
15) L. Ostrach and J. Head, LPSC 38, 1100, 2007; 16) J. Dickson et
al. Vernadsky-Brown Micro 46 , 2007; 17) E. Hauber et al.,
EMSEC:MEE , 2007; 18) J. Head and D. Marchant, Geology, 31,
641, 2003; 19) F. Forget et al. Science, 311, 368, 2006; 20) J. Head
et al., GRL, 31, L10701, 2004; 21) J. Head and S. Pratt, JGR, 106,
12275, 2001.
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EVIDENCE FOR WATER BY MARS ODYSSEY IS COMPATIBLE WITH A BIOGENIC DDS-
FORMATION PROCESS. T. Gánti (1), A. Horváth (2,3), Sz. Bérczi (4), A. Gesztesi (3), E. Szathmáry (1,5)
(1) Collegium Budapest (Institute for Advanced Study), 2 Szentháromság, H-1014 Budapest, Hungary, ; (2) Konkoly Observa-tory, H-1525 Budapest Pf. 67, Hungary; (3) Budapest Planetarium of Society for Dissemination of Scientific Knowledge, H-1476Budapest Pf. 47, Hungary, ( [email protected]) ; (4) Eötvös University, Dept. G. Physics, Cosmic Materials Space Research
Group, H-1117 Budapest, Pázmány 1/a. Hungary, ( [email protected]); (5) Eötvös University, Dept. of Plant Taxonomyand Ecology, H-1117 Budapest, Pázmány 1/a. Hungary, ([email protected]);
Abstract: The neutron measurements of the Mars Odysseyspacecraft provided evidence for some form of water in the
upper layer of the surface of the Southern Polar Region of Mars. This is compatible with the prediction of a presumablyexisting supply of water in our model of DDS formation,where the MSOs (Mars Surface Organisms) are covered by
seasonal water ice. Hypothetical MSOs are thought to meltthe ice above them, which initiates the characteristic
morphogenesis of DDSs [1, 2 and 3]. Concomitant melting
of water frost in the uppermost layer of the soil may contrib-ute to the observed traces of liquid water.
Evidence for water by the Mars Odyssey on the South
Pole: MGS MOLA data indicated that the southern freshfrost cover is 0.1--1 meter thick [4]. (These are averagevalues, from which significant deviations may occur locally.)However, the frost cover surely consists of three compo-
nents: frozen carbon dioxide, carbon dioxide clathrate and
water ice [5, 6]. Unfortunately, till today we do not have anydata about the depth of these layers.
Fig. 1. Evidence for water by the High-Energy Neutron Detector (HEND) of the Mars Odyssey spacecraft from February to April of
2002 in the South Polar Region of Mars (in the summer), and the
sites of dark dune spots (red circles).
Data provided by the Gamma Ray Sensor on the MarsOdyssey spacecraft by the American Neutron Spectrometer
(NS) and the Russian High Energy Neutron Detector (HEND) indicated deficits of high-energy neutrons in south-ern highlands of Mars. (These deficits indicate that hydrogenis concentrated in the subsurface.)
Model calculations suggested that the best possible fit to
the data was a water-ice rich layer with at least tens of centi-meters in thickness. In this case the subsurface material has a
water ice concentration of 60% by volume [7, 8 and 9].This prevalence of water ranges from the South Pole up
to 60ºS, surprisingly coinciding with the region of the DDSs(Fig. 1). From this data we may deduce that water in some
form is relatively abundant in this region of the DDSs.
DDS formation process: Dark Dune Spots (DDSs) and their clusters (Fig. 2) are interesting objects with seasonal frost
cover transformational dynamics. They were observed on theMars Orbiter Camera (MOC) narrow angle images of the
Mars Global Surveyor (MGS) spacecraft [10] from the year
1998 to 2002.The frost cover appears during autumn and graduallythickens in winter until formation of DDSs transforms thislayer, beginning in late winter. On surfaces other than the dark dunes the frost disappears during spring, but on the latter the
frost persists until late spring or even early summer [11].
Fig. 2. Example for the characteristic features of a dark dune spot
(DDS) field ( b ) in the crater Chamberlin ( a ). Sun in the figures
illuminates from upper left, north is up.
In the earlier detailed analyses of several thousand dark dune spots on the Southern Polar Region of Mars from MGSMOC images we could determine the shape, the pattern and
the seasonal/annual dynamics of these spots [2, 3, 12, 13]. We studied the transformational process according to a
hypothesis which describes the observations on dark dune
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Gánti, Horváth, Bérczi, Gesztesi, Szathmáry:EVIDENCE FOR WATER BY MARS ODYSSEY IS COMPATIBLE WITH A BIOGENIC DDS-FORMATION PROCESS
spots morphology, the patterns of DDSs, the formation andtransformation of DDSs by using the following components:
the Dark Dune fields’ material and its frost cover , the putativeMartian Surface Organisms (MSOs) and their supply of water (embedded in the hypothetically porous dark dune soil).
Individual DDSs exhibit a characteristic spot structure
with an inner dark region, transitional gray peripheral ring
and the frost cover surrounding the structure (Fig. 3a) [1, 2, 3and 13]. On slopes the spots are elongated downwards: they
occur in ellipsoidal or fan-shaped forms and sometimes
streams flow out from these spots (Fig. 3b) [1, 2, 3 and 13].We also observed that DDSs are shallow crater-like holes inthe frosted layer and that the DDS formation process is trig-
gered from the bottom of the frost (Fig. 3c,d) [3, 12]. We
also observed seasonal and annual variation and recurrence
of DDS patterns, too (Fig. 4a-d and 4e,f ) [3, 12]. The shape,location, development and other features of the DDSs
prompted us to suggest that some fluid phase must be in-
voked in their explanation, which under the given circum-stances cannot be anything else but liquid water.
Fig. 3 Dark dune spot (DDS) characteristic features: inner ring
structures ( a ), fan and flow-shaped forms ( b ), shallow holes-form
( c ). Sun in all figures illuminates from upper left north is up.
Biological interpretation of DDS-formation by four
agents: We interpret the sequence of DDS formation as anaccelerated process of sublimation combined with melting of water and some kind of biological activity of putative Mars
Surface Organisms (MSOs) acting on, or in, the material of
the dark dunes [1, 2, 3, 12 and 13].
Fig. 4 Dark dune spots seasonal ( a-d ) and annual ( e, f ) dynamics.
Sun in all figures illuminates from upper left north is up.
We sketched the following life cycle of MSOs: in winter the first rays of sunlight activate the MSOs, they start towarm up and melt the H
2O ice around them, while above
them sublimation of CO2 on the top of the frost is acceler-ated. Later MSOs begin to grow and reproduce themselves inthe water melted by them. Complete defrosting of the water
ice cover stops shielding the MSOs and water immediatelyevaporates, too, on this unprotected region, hence the lifeconditions of MSOs cease and they desiccate.
All these events happen in the upper layer of the dark dune soil. In our model we suggest that the four componentsact together in a life-desiccation cycle, alternately followingeach other seasonally and annually. Therefore, despite theadverse conditions, the hypothetical Martian Surface Organ-
isms could dwell below the surface ice, in the upper water-rich layers of the dark dune field. Melting of water ice in theuppermost layer of the dune subsurface, triggered by the life
activity of MSOs above, may contribute to the observedtraces of liquid water associated with the DDSs.
Conclusions: Neutron measurements of Mars Odyssey
(HEND, NS) observed the presence of water ice, which wasabundant in the upper 2-meter thick layer of the SouthernPolar Region between 90°S--60°S latitudes [7, 8 and 9]. Thisobservation is compatible with prediction for the existence of
seasonally melting water on the basis of the MSO hypothesis,and that the DDS phenomenon is evidence for present-daylife on Mars.
Acknowledgments: Authors thank for the use of MGS MOC images of NASA
JPL and Malin Space Science Systems (http://www.msss.com/moc_gallery/).
References: [1] Horváth et al. 2001, Probable evidences of recent biological
activity on Mars: Appearance and growing of dark dune spots in the South
Polar Region, LPS XXXII, # 1543. [2] Gánti et al. 2002, LPS XXXIII , #1221.
[3] Gánti, Horváth et al. 2003, Dark Dune Spots: Possible Biomarkers on
Mars? OLEB in print. [4] Smith et al. 2001, Science 294, 2141. [5] Hoffman
2000, Icarus 146, 326. [6] Carr 1996, Oxford Univ. Press. [7] Boynton et al.2002, Science 297, 81-85. [8] Feldman et al. 2002, Science 297, 75-78. [9]
Mitrofanov et al.2002, Science 297, 78-81. [10] Albee et al. 2001, J. Geophys.
Res.106 E10, 23291. [11] Malin M., Edgett K. 2001, J. Geophys. Res 106.
23429. [12] Horváth et al. 2002a,b, LPS XXXIII, # 1108, 1109. [13] Horváth et
al. 2002c, Antarctic Meteorites XXVII, 37. Tokyo.
Lunar and Planetary Science XXXIV (2003) 1134.pdf
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
EXOMARS ENTRY AND DESCENT SCIENCE
F. Ferri1, A. J. Ball
2, G. Colombatti
1, A. Aboudan
1, F. Angrilli
1, I. Müller-Wodarg
3, B. Hathi
2, M. R. Leese
2,
S. R. Lewis2, J. C. Zarnecki
2and the EDLS Science Team.
1CISAS “G. Colombo”, University of Padova, Via Venezia 15, 35131 Padova, Italy.
2Centre for Earth, Planetary, Space and Astronomical Research, The Open University, Walton Hall, Milton
Keynes MK7 6AA, UK.3Space and Atmospheric Physics Group, Imperial College London, Prince Consort Rd, London SW7 2BW, UK.
Email: [email protected] , [email protected] .
The entry, descent and landing of ExoMars offer
a rare (once-per-mission) opportunity to perform in
situ investigation of the martian environment over a
wide altitude range. We present an initial
assessment of the atmospheric science that can be
performed using sensors of the Entry, Descent and
Landing System (EDLS), over and above the
expected engineering information. This is intended
to help fulfill the concept of an Atmospheric
Parameters Package (APP), as mentioned in the ExoMars draft Science Management Plan [ESA,
2005].
Mars’ atmosphere is highly variable in time and
space, due to phenomena including inertio-gravity
waves, thermal tide effects, dust, solar wind
conditions, and diurnal, seasonal and topographic
effects. Atmospheric profile measurements, drawing
on heritage from the Huygens Atmospheric
Structure Instrument (HASI), which encountered
Titan’s atmosphere in 2005 [1], should allow us to
address questions of the martian atmosphere’s
structure, dynamics and variability.
Figure 1. ExoMars EDLS sequence
By careful definition of EDLS measurements to
yield science as well as a successful landing, we aimto obtain continuous atmospheric density,
temperature and pressure profiles over the widest
ever altitude range, with the highest sensitivity and
spatial resolution.
Extrapolation to the ExoMars case of the flight
performance of the HASI entry accelerometry
experiment is encouraging.
Up to now, only three high vertical resolution
and high accuracy vertical profiles of density,
pressure and temperature of the martian atmosphere
have been derived from in situ measurements
performed by Viking 1 and 2 in day-time [2] and by Mars Pathfinder in night-time [3, 4]. Two more
vertical profiles have been retrieved from the
deceleration curves and aeroshell drag properties of
the two Mars Exploration Rovers (MER) during
atmospheric entry [5], but with a much lower
accuracy.
Such profiles are vital for testing of atmospheric
models used in numerous studies of atmospheric
variability, on a range of temporal and spatial scales,
as well as for the practical issue of reaching the
martian surface reliably [6].
New data from different site, season and time
period are essential to investigate the thermalbalance of the surface and atmosphere of Mars,
diurnal variations in the depth of the planetary
boundary layer and the effects of these processes on
the martian general circulation.
A better understanding of the martian
environment and meteorology is also essential for
refining and constraining landing techniques at Mars
and to evaluate the possible hazardous to machines
and humans in view of future Martian explorations.
As the ExoMars project definition proceeds, the
entry, descent and landing sequence may offer
further science opportunities. We would be
interested in exploring these and welcome additional
members to the consortium. The joint team co-
ordinators are Francesca Ferri (Univ. Padova, Italy)
and Andrew Ball (Open University, UK).
References:
[1] Fulchignoni, M. et al. (2005), Nature 438(7069), 785-
791.
[2] A. Seiff, D.B. Kirk, (1977) J. Geophys. Res 82,. 4364-
4378,.
[3] Schofield, T., et al. (1997) Science 278, 1752-1758
[4] Magalhães, J.A., J.T. Schofield, A. Seiff, (1999) J.
Geophys. Res.104, 8943-8945.
[5] Withers, P. and M. D. Smith (2006) Icarus 185, 133-
142.[6] Montabone, L. et al. (2006) Geophys. Res. Lett.
33(L19202).
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THE EXO-MARS EXPERIMENT MINIHUM A. Koncz1, V. Schwanke, D. T. F. Möhlmann, R. Wernecke,
A. Lorek 1DLR Inst. of Planetary Research, 12489 Berlin, Germany.
2Dr. Wernecke & Partner, 14480 Potsdam,
Germany. [email protected]
The water content of soils significantly influencestheir chemical, physical and biological properties. In
respect of Mars the thin layer of the upper
millimetres of the Martian surface are of particular
interest since this soil interacts directly with the
diurnally varying atmospheric humidity, which can
reach saturation during night and early morning.
Adsorption/desorption of water in the soil and
freezing of it can be a consequence. Therefore, near-
surface measurements of the atmospheric content of
water vapour will allow to investigate the
interaction between the atmosphere and the
adsorbed water, which is deposited in the upper
soillayer. This should be an important point inrespect of exobiology and the exploration of Mars.
The Institute of Planetary Research at the German
Aerospace Center and the SMB Dr. Wernecke &
Partner are jointly developing a humidity sensor
system in preparation for the ExoMars mission. Its
goal is to obtain first in-situ measurements of
diurnal and seasonal variations of the near-surface
atmospheric water vapour content. By consequent
miniaturization the MiniHUM Team was able to
reduce size and weight of the coulometric sensor in
such a way that the system, weighing 180g, will be
one of the smallest instruments on board of the
ExoMars mission. The instrument itself consists of two different units: HUM and ASS. The humidity
sensor HUM is a combined unit of the coulometric
sensor (QSE) and a capacitive (CPS) humidity
sensor with an integrated thermocouple. The unit
will be placed inside a small sensor housing (55 x
48 x 9 mm) on the outer solar panel of the lander
structure. For measuring the humidity under Martian
conditions the principles of coulometric and
capacitive measurement are the most appropriate.
The coulometric principle is based on the ability of Diphosphorpentoxid (P2O5) to adsorb environmental
water vapour almost completely, generating finally
HPO3 in solution. In case of a sufficiently applied
DC voltage this leads to a resulting electric current
due to charge separation in the solution. The
resulting current is directly related to the amount of
adsorbed water, as described by Faraday´s law. This
offers the possibility to quantitatively measure
absolute humidity in the dew-point range between –
90° C and -40°C. The capacitive humidity sensor
uses the humidity dependence of some polymeric
dielectrics to measure the relative humidity content
of environmental gas from 5% r.h to 80% r.h. Byusing the two different principles of measurement it
is possible to ensure a wide working range but also
redundancy and a cross-reference for each sensor.
The Atmosphere Saturation Sensor (ASS) will be
mounted externally and consists of a high sensitive
resistant thermometer for independent determination
of the frost point temperature. The measurement of
the frost point temperature will give an additional
and independent way to determine absolute
humidity, specifically at that point of phase
transition (cf. Ryan and Sharman, 1981). This
information can also be used for calibrationpurposes of the humidity sensors.
References:
Ryan, J.A.; Sharman, R.D. (1981), H2O frost point
detection on Mars, Journal of Geophysical Research 86
(1981), S. 503-511
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
EXPECTED MAGNETIC SIGNATURE OF DEMAGNETIZED IMPACT CRATERS TO TIME THE
DYNAMO SHUTDOWN ON MARS. B. Langlais1, E. Thébault2, Y. Quesnel3, C. Sotin4, and F. Leblanc5.1Laboratoire de Planétologie et Géodynamique de Nantes, CNRS/Université de Nantes, Nantes, 44322, France.2Institut de Physique du Globe de Paris, France. 3GFZ Potsdam, Germany. 4Jet Propulsion Laboratory and
California Institute of Technology, Pasadena, CA, U.S.A. 5Service d’Aéronomie du CNRS/IPSL, Université
P&M Curie, Paris, France. [email protected]
Introduction: Mars Global Surveyor (1996-
2006) detected strong and localized magnetic field
on Mars [1]. They are the relics of an ancient, likely
planetary scale, dynamic magnetic field. The
magnetic field was frozen by the magnetic minerals
of the lithosphere, and its remanent trace has not
evolved since, unless re-heated, displaced or
shocked, resulting in a de- and eventually re-
magnetization. In this paper, we study the
demagnetization associated with an impact crater,
and predict its magnetic signature at spacecraft
altitude. By comparing this predictive signal to low-latitude spacecraft magnetic field measurements, it
should be possible to discriminate between impact
craters that did not affect the magnetic properties of
the lithosphere, and those who did (occurring after
the Martian dynamo shutdown).
Demagnetization and impact craters: Impacts
are associated with thermal and shock phenomena.
But the first demagnetization process associated
with impacts is material excavation. In the following
we focus on this signature. Impact craters are
described by several parameters. The final,
measurable, rim-to-rim crater diameter Dt
is
empirically expressed as a function [2] of the
transient one Dr. The excavation depth d
exrepresents
the maximum depth at which the material was
excavated. The excavation depth-to-diameter ratio is
approximated by dex /D
ex= d
ex /D
t~ 1/10.
Pre-impact material: We assume themagnetization to lie in the upper 60 km of the
Martian litho-sphere. The thermo remanent
magnetization was acquired while cooling in the
presence of a global, axial dipolar magnetic field.
Magnetization direction vary with the location of
the impact on the globe. The mean magnetization is
set to 1 A/m, comparable with the mean
magnetization of existing models [3]. The excavated
area is simulated by a paraboloid of revolution, with
a circular surface section and a maximum depth set
to 10% of its diameter. Crater diameters range
between 100 and 400 km, with a 50-km increment.
Results: The three components of the magnetic
field are predicted at altitudes ranging from 100 to
400 km. The resulting magnetic field signature is
more or less symmetric, depending on the location
of the crater. The largest field is observed when the
impact is emplaced above the magnetic pole. At
200-km altitude, a 200-km wide demagnetized
impact crater results in a ~8 nT magnetic field. Such
a transient crater corresponds to a rim-to-rim
diameter close to 350 km.
The larger the crater is, and the lower the
measurement is made, the larger the magnetic field
is. A 100-km diameter transient crater has a
magnetic signature of about 3 nT at 150-km altitude,
while a 200-km diameter crater is associated with a
1 nT signal at 400-km altitude.
Discussion: Large and small impact craters may
affect the magnetic properties of the lithosphere.
Individual craters located in a homogenously
magnetized layer have magnetic signatures that aremeasurable at spacecraft altitude (150 km and
above). A single crater is associated to a local
increase of the magnetic field. The case for Mars is
much more complex, as multiple overlaying craters
are emplaced within a non-homogeneously
magnetized lithosphere.
Finding one-to-one correlations is not easy, since
many parameters influence the magnetic signature.
But it should be possible to determine the
characteristics of demagnetized areas [4]. Some very
large craters (> 300 km), such as Newton and
Copernicus, may be partially demagnetized. Others,
such as Daedalia, do not show any correlation withthe magnetic field. The two first craters are
consequently thought to have taken place after the
dynamo shutdown. Characterization of smaller
craters is more difficult, as the existing
measurements are either too high or too sparse.
Conclusion – the case for MEMO: New
measurements are eagerly awaited for. These new
measurements should have a better accuracy (~ 0.5
nT), a much better geographical coverage at low
altitude (below 250 km). This is one of the main
scientific objectives of Mars Escape and Magnetic
Orbiter (MEMO)[5]. Such an orbiter would
simultaneously monitor the atmospheric dynamicsat Mars and the interactions between the
lithospheric magnetic field and the solar wind. The
low altitude orbit of MEMO would allow the
magnetization of craters as small as 100-km
diameter to be characterized. There are more than
250 such craters. These craters will help to better
estimate the dynamo shutdown on Mars, which will
then bring new constraints on the evolution of Mars’
interior and of its atmosphere.
References: [1] Acuña M.H. et al. (1998)
Science, 279. [2] Croft S.K. (1985) Proc. LPSC XV ..
[3] Langlais B. et al. (2004) J. Geophys. Res., 109.
[4] Thébault E. et al. (2006), EPSC , Abstract #244.
[5] Leblanc, F., et al. (2007), EMSEC conf.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
EXPLORATION AND CHARACTERIZATION OF A NOVEL THERMOPHILIC
BACTERIUM, KURTHIA SP. A. Kumar, R. Kumar. Institute of Genomics and Integrative Biology, Mall road,
Delhi, India. [email protected]
Thermophilic microorganisms have been of
great scientific interest for several decades,principally in regard to their
biotechnological potential and also of the
possibility of their existence in extreme
exobiological econiches1.Present work deals
with the isolation of an extremophile, its
identification and application in
bioremediation. For isolation, water sample
of high alkalinity and high temperature
(55oC) was taken from the textile industrial
premises. On the basis of high temperature
tolerance and high pH (12.0), a gram
positive and aerobic bacterium was
screened. This bacterium was identified asKurthia sp., which is capable to hydrolyze
the starch. This bacterium was exploited to
neutralize the hot and alkaline industrial
wastewater from pH 12.0 to pH 7.5.
Alkaline bacillus medium (ABM) was
selected as the suitable medium to grow this
bacterium. For the neutralization of hot andalkaline wastewater, ten hours grown
culture was centrifuged and the pellet was
added to hot (55oC) wastewater of pH 12.0.
Lowering of pH from 10.0 to 7.5 using this
bacterium could be achieved in a period of
one hour. This kind of bacteria, which are
capable to grow in extreme conditions, can
provide some insight to explore the life in
other planets.
References:
1. Fujiwara, N. and K. Yamamoto (1987), Decompositionof gelatin layers on X-ray film by the alkaline protease
from Bacillus sp. B21. J Ferment Technol, 65:531/534
.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Exploring Mars with martian meteorites: casting light on primary, secondary and tertiary processes. Monica M.
Grady1,2
and I. P. Wright1.
1PSSRI, The Open University, Milton Keynes UK.
2Dept. Mineralogy, The Natural
History Museum, London, UK. [email protected]
There are currently about 60 meteorites from
Mars (38 distinct samples), all of which are igneousrocks, sub-divided into four groups and eight sub-
groups on the basis of mineralogy and geochemistry
[1]. Each group represents rocks that formed in
different locations at or below the martian surface.
They cannot all have come from a single impact
event - at least three craters, with minimum
diameters of about 12 km, are required to produce
the variety of martian meteorite types [2].
Measurements made on martian meteorites
complement data from spacecraft exploration of
Mars, and until we have a mission that returns
material directly from the planet, these meteorites
are the only physical objects that we have to help usunderstand martian geology and processes on Mars.
Shergottites, the most numerous of all martian
meteorites, are the youngest rocks (ages < 600 Ma;
ref. 2). Nakhlites crystallised 1.3 billion years ago
[2], i.e. at the boundary between the early- and mid-
Amazonian epochs [3], and emanate from different
depths from a sill or dyke emplaced close to the
martian surface [4]. They contain microscopic
evidence of secondary minerals and features
resulting from interaction with water, including
iron-rich carbonates, sulphates, halite and clay
minerals [5]. The iron-rich carbonates are intimately
mixed with the clay minerals, implying that they
formed together, or at least in the same fluid
alteration event [6]. It has proved difficult to
measure the age of the clay minerals, because of
their fine-grained nature, but on the basis of K-Ar
data, they are thought to be up to 700 Ma old [7, 8].
There are no reported ages for carbonates in
nakhlites, but they must be younger than their 1.3
Ga crystallisation age, and the < 700Ma age for the
clay minerals is more likely to be an age for a clay-
carbonate mixture. The secondary minerals form a
sequence that has been interpreted as a mineral
assemblage produced by progressive deposition
from an evaporating brine [6]; if Mars, in the past,
had been colder and drier than previously thought,
then this alteration scenario will have to be re-
evaluated. Chassignites have the same
crystallisation age as the nakhlites, but show very
few aqueous alteration features. They are presumed
to come from a deeper burial depth than nakhlites.
ALH 84001 has a mineralogy, petrology and thermal
history that sets it apart from the other martian
meteorites [9]; it is the oldest of all (early Noachian;
crystallisation age ~ 4.5 Ga; ref. 2). Although it
contains an abundance of carbonates (~ 1 vol. %),
there are few other signatures of alteration, implying
an unusual history of secondary alteration [10].Age-dating of the carbonates in ALH 84001 has
shown them to be ~ 3.9 Ga old [11, 12], younger
than the formation age of the meteorite [2], but close
to its shock age [13].
There are three phases of information we can
gain from martian meteorites. The first is
information about primary magmatism on Mars.
Because the meteorites are igneous rocks, study of
the composition and mineralogy of these rocks helps
us to understand the temperature and mode of
volcanism on Mars. Secondary alteration products
in these meteorites helps us to learn about fluid flow
on or near the surface of Mars. Particularly in ALH84001 and in the nakhlites the complex assemblages
of secondary minerals shed light on the temperature
and salinity of the water that once flowed across
Mars’ surface. The zoned nature of some of the
minerals tells us how fluid composition has
changed, either in terms of temperature or in terms
of the salts dissolved in the fluid. The restricted
nature of the alteration assemblages indicate that
fluid flow was limited, perhaps confined to an
evaporating basin rather than to a river or a stream.
The third phase of information that we can learn
from martian meteorites is that which is associated
with shock. The shergottites and the chassignites areshocked meteorites containing patches of diaplectic
glass (maskelynite) formed by the conversion of
plagioclase during shock melting and quenching.
Trapped within these melt pockets are martian
atmospheric gases. By analysing martian meteorites
of different ages, it is possible that we might be
looking at samples of Mars’ atmosphere trapped at
different times. In this way, it is possible that we
might be able to trace the evolution of Mars’s
atmosphere by looking at these pockets of gas.
References: [3] Hartmann W. K. & Neukum G. (2001)
Sp. Sci. Rev. 96, 165-194; [5] Bridges J. C. et al. (2001)
Sp. Sci. Rev. 96, 365-392; [1] Bridges & Warren (2006);
[2] Nyquist et al. (2001) Sp. Sci. Rev. 96, 263-292; [4]
Treiman A. H. (1986) GCA 50, 1061-1070; [6] Bridges J.
C. & Grady M. M. (2000) EPSL 176, 267-279; [7]
Swindle T. D. et al. (2000) MAPS 35, 107-115; [8] Shih
C.-Y. et al. (1998)LPS XXIX , Abst. No. 1145; [9] Treiman
A. H. (1998) MAPS 33, 753-764; [10] Warren P. H.
(1998) JGR 103, 16759-16773; [11] Borg L. E. et al.
(1998) In: Where do we stand & where are we going?
Abst. No. 7030; [12] Borg L. E. et al. (1999) LPS XXX ,
Abst. No. 1430; [13] Turner G. et al. GCA 61, 3835-3850,
1997
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
FIRST ESTIMATION OF TOTAL ELECTRON CONTENT OF MARS IONOSPHERE A. Safaeinili1,
W. Kofman2, J. Mouginot
2J. Plaut
1, Giovanni Picardi
3,1Jet Propulsion Laboratory, Pasadena, 91109, USA.
2
Labratoire de Planetologie de Grenoble, CNRS/UJF, France,3Dipartimento INFOCOM, University of Rome
``La Sapienza'', Rome 00184, Italy. [email protected] ,
Introduction: We present results on estimationof Mars ionosphere’s total electron content (TEC)
using a novel transmission mode technique. We
have exploited the frequency dependent phase
distortion in the surface echo signal of MARSIS
radar to estimate the TEC of the ionosphere. This
technique can provide variation of the TEC at a
spatial resolution of 5-10 km which equivalent to
one radar footprint of MARSIS. So far we have
obtained more then 0.5 million independent
estimates of TEC over Mars under varied sun
elevation angles. We have compiled these results
to provide a more accurate average model of the
Mars ionosphere. We have also observed intriguingconnection between nighttime TEC enhancements
and the crustal magnetic field of Mars.
Radar Observations: MARSIS echoes can be
obtained at any of the four MARSIS subsurface
sounding bands with center frequencies of 1.8,
3.0,4.0 and 5.0 MHz with a one MHz bandwidth.
The roundtrip propagation of radar pulse through
the ionosphere causes frequency dependent slowing
of the wave. This dispersion encodes the
information about the ionosphere column between
the spacecraft and the ground, as shown in Fig. 1.
This information is used to calculate the TEC (asshown in Fig. 2) and the Chapman ionospheric
height scale. The distributions scale height and
peak plasma frequency are bi-modal, which seems
to be related to the local time of the observations.
Separating the data by the local time of the
terminator crossing, i.e. whether the pass included
the sunset or sunrise terminator (the satellite moving
from small SZA to large or inverse), we get two
distinct populations. The sunrise orbits tend to
have a smaller scale height of 11.5 km, and the
sunset orbits have a scale height of 15 km.
Figure 1. Mars ionosphere distorts the radar signal: Upper
panel shows the radar signal distortion as the spacecraft
travels from nightside to the dayside, Lower panel shows
the corrected radargram
Figure 2. The ionospheric correction provides an estimate
for the TEC. This particular example is derived for the
example shown in Figure 1.
Similarly, the equivalent extrapolated N0 is higher
for sunrise orbits at ~ 2.1x1011
m-3
and is lower for
the sunset orbits at about 1.3x1011
m-3
. The sunset
and sunrise parameters can be considered as
bracketing the range of actual values during the
diurnal cycle.
Our data also indicates the presence of a high
concentration of electrons in some regions at the
night side of Mars. These regions are very abrupt
and isolated and seem to correspond with regions
where Mars crustal magnetic field is open, which in
general corresponds to areas with vertical magnetic
field vector. The regions with enhanced plasma
have been observed by other workers [2,3]. As the
Mars Express pericenter passes drift more into the
night, we will be able to greatly increase our
coverage of Mars ionosphere’s nightside behaviour.
Acknowledgments: MARSIS is a joint project of
ASI and NASA onboard ESA’s Mars Express
Spacecraft. Some of the work described herein was
performed at the Jet Propulsion Laboratory under
contract with NASA.
References: [1] Safaeinili, A. et al. (2003), Planetary and
Space Science, 51, pp. 505-515, 2003. [2] Nielsen, E., et
al., Vertical sheets of dense plasma in the topside Martian
ionosphere, J. Geophys. Res., 112, 2007. [3] Duru. F., et
al. , “magnetically controlled structures in the ionosphere
of Mars”, JGR, Vol. 3, pp. 12204, 2006.
Seventh International Conference on Mars 3206.pdf
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European Space AgencyEuropean Mars Science and Exploration Conference: Mars Express & ExoMarsESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
FIRST MODELLING OF THE MARS IONOSPHERE WITH THE EUROPEAN MARS GCM
G. Gilli1, F. González-Galindo
2, M. A. López-Valverde
1, F. Forget
2.
1Instituto de Astrofísica de Andalucía–CSIC, Apdo. 3004, Granada 18008, Spain.
2Laboratoire de Météorologie
Dynamique, IPSL, Université Pierre & Marie Curie, BP99, 4 place Jussieau, 75252 Paris, CEDEX 05, [email protected]
IntroductionThe purpose of this investigation is the design,
implementation and testing of a model of theMartian ionosphere into the thermosphericextension of the European Mars General CirculationModel (EMGCM) [1]. We have two main scientific
objectives for adding this new module to theEMGCM. First, we want to study the coupling of the plasma to the neutral chemistry and dynamics of the Martian thermosphere, as currently simulated bythe EMGCM, so that the possible effects on themodel's current energy budget and dynamics at
those altitudes can be evaluated. And second, thisnew tool shall make the EMGCM suitable foranalysis of the recent ionospheric measurements byMars Express and other missions [2,3]. In addition,we will have new validation capabilities for theEMGCM at those altitudes, like comparison withthe NCAR-UMI Mars Thermospheric GCM [4,5].The strategy is, first, to develop and test a simple
ionospheric model in a 1-D model of the Martianupper atmosphere [6]; second, to design fastversions for implementation into the currentEMGCM, similar to those already used for theneutral chemistry [7]; and third, to test thenumerical behavior and physical coupling of theionosphere within the full EMGCM.
Model characteristics Our 1-D Martian model and the thermospheric
version of the EMGCM extends from the surface upto above 200 km, and this is the nominal altitude of our upper boundary.
A number of ionospheric models of differing
sophistication have been developed/applied to Marsin the past [8,9]. Our ionospheric model includes inthis first version only those reactions necessary for acorrect simulation of the nature and variability of the main ionospheric electron peak, and thoserequired for internal consistency and mass
conservation within the model. Additionalapproximations include maintenance of globalneutrality of the plasma, and neglect of reactionswith secondary electrons, magnetic disturbances,and pick-up ions from the solar wind.With the aid of the 1-D model we evaluated theimportance of a large number of reactions of theneutrals with the most important ions, O2+, O+,
CO2+, CO+, N+, NO+, N2+, H+, in addition to therequired photoionizations and dissociativerecombinations, and we ended up with a set of 37selected reactions, which include the more relevantionization channels. For the dependence of rate
constants on electron temperature, we simplyassumed a given reference value [10].
GCM implementation
Examination of chemical productions, losses and
lifetimes of all the species included in the 1-Dmodel, allowed us to confirm the validity of thephotochemical equilibrium approximation for all theions and the small impact on the neutrals, as well asto design a fast calculation module with minorfurther assumptions, which is suitable for the GCM.The study included numerical stability and time-
marching consistency tests, and analysis of arbitraryinitial conditions and the behavior during day-nighttransitions.
We will present first simulations of the
ionospheric module within the 1-D and the GCM,devoted to investigate the natural variability of thealtitude and magnitude of the Martian ionospheric
peak.
References: [1] Angelats i Coll et al., GRL, 32, L04201(2005); [2] Pätzold et al., Science, 310, pp.837-839(2005); [3] Fox and Yeager, JGR, 111, A10309 (2006);[4] Bougher et al., JGR, 109, E03010 (2004); [5]Gonzalez-Galindo et al., this issue; [6] Lopez-Valverde etal., II Workshop Mars Atmosphere Modelling andObservations, Granada (2006); [7] Gonzalez-Galindo etal., JGR, 110, E09008 (2005); [8] Krasnopolsky V.A.JGR, 107, E12 (2002) [9] Moffat T., PhD Thesis, UCL
Londond (2005), [10] Hanson et al. JGR, 82, pp.4351-4363 (1977)
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
GENERATION OF HIGH-RESOLUTION DIGITAL TERRAIN MODELS AND ORTHO-IMAGE
MOSAICS FROM OLYMPUS MONS, MARS, ON THE BASIS OF MARS-EXPRESS/HRSC DATA A.
Dumke1, M. Spiegel2, R. Schmidt3, G. Neukum1. 1Institute of Geoscience, Freie Universität Berlin, Malteserstr.
74-100, 12249 Berlin, Germany. 2Photogrammetry and Remote Sensing, Technische Universität München,
Arcisstr. 21, 80333 Muenchen, Germany. 3Institute of Photogrammetry and GeoInformation (IPI), Leibniz
Universität Hannover, Nienburger Str. 1, 30167 Hannover, Germany. [email protected]
Introduction: Since December 2003, the European
Space Agency’s (ESA) Mars Express (MEX) has
been orbiting Mars. The High Resolution Stereo
Camera (HRSC), as part of the scientific
experiments onboard MEX, is a pushbroom stereo
colour scanning instrument with nine CCD line
detectors and 5184 sensor elements per line. It
consists of five lines with panchromatic filters and
four lines with red, green, blue and infrared filters
[1]. Until now, the HRSC has covered an area of
approx. 100 million km2, up to a spatial resolutionof 10-20 m per pixel. These are well suited to
derive DTMs, ortho-image mosaics and additional
higher-level 3D data products.
Geoscientific studies can be carried out in single-
orbit image data, but in order to obtain a more
comprehensive view of regional processes on Mars,
images as well as topographic data have to be
mosaicked photogrammetrically. One of the primetargets for such studies is the Olympus Mons
volcano.
Methods: Olympus Mons was imaged by theHRSC instrument during 18 MEX orbits. 16 HRSC
orbit strips show good image qualities and were
used to generate a DTM mosaic as well as a gapless
ortho-image mosaic covering an area of approx.
600,000 km2. The nadir ground resolution is in therange of 12 m to 40 m per pixel.
Derivation of DTMs and ortho-image mosaics are
basically performed using software developed at
the German Aerospace Center (DLR), Berlin and
based on the VICAR tools developed at JPL. The
standard processing workflow is described in detail
in [2,3]. The main processing tasks are (a) pre-rectification by using the global MOLA-based
DTM, a least-squares area-based matching between
nadir and other channels (stereo and photometry)
and (c) DTM raster generation. Parameters for the
derivation of preliminary DTMs are individually
adapted to the quality of orbit data and imagequality. Additionally, iterative image filtering is
applied in order to improve the image matching
process by increasing the amount and quality of
object points.
Apart from the DTM quality, image mosaicking
also depends on the quality of exterior orientation
data and in order to generate high resolution DTMs
and ortho-images, these data have to be corrected.
For this purpose, new exterior orientation data
based on tie point matching and bundle adjustment
provided by the Leibniz Universität Hannover and
Technische Universität München have been used
[4,5]. This allows us to adopt HRSC-derived data
to the global Mars-reference system as defined by
MOLA. The new orientation data refinements have
been applied for individual strips thus far.
Additionally, there are bundle-block adjustments
for five combined orbits covering the eastern partof Olympus Mons.
Results: DTM derivation using exterior orientationdata that were adjusted in a strip could be used for
ortho-image mosaics and DTMs, and provided
good results. As expected, the five orbits that were
adjusted in a block have a higher accuracy when
compared to the orbits adjusted in a single strip.
The panchromatic HRSC ortho-image mosaic hasbeen generated successfully and methods described
above will now be extended to adjacent nadir strips
and the colour channels
A topographic map based on HRSC derived image
data and DTM data will be presented in a poster
(Fig. 1).
Figure 1. Ortho-image mosaic with superimposed colour
coded elevation
References: [1] Neukum et al. (2004), ESA SP-1240,
17-35, [2] Scholten et al. (2005), PFG 5, 365-372, [3]
Gwinner et al. (2005), PFG 5, 387-394, [4] Schmidt et al.(2005), PFG 5, 373-379, [5] Spiegel et al. (2007), Lunar
and Planetary Science XXXVII, Abstract #1608
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
MATHEMATICAL ASPECTS ASSOCIATED TO THE GENERATION OF IMPACT CRATERS ON
MARS J. C. Echaurren1.
1Codelco Chile Chuquicamata, North Division, Pasaje Lince 976, Calama, Chile.
Introduction: The aim of this work is to
generate mathematical models associated to thecreation of impact craters, through the use of the
physical properties relationed with both asteroids
and target rocks. Will be showed here, the principal
expressions [1,2] involved in the calculations for the
obtention of impact conditions on craters generated
by asteroids or comets on Mars. These equations are
basically of empirical order and numerical
adjustment, for the considered impact conditions.
Mathematical Formulation: The models used
in this work are relationed with mass distributions of
polynomial order [1], associated to asteroids in
movement (dynamical conditions), which aredescribed as,
A2m24 – A2m1
4 + 2A(A – gP)m1 m23 + 2A(gP – A) m1
3
m2 = 0 ,
where,
A = (V ATM / 2)2 = constant of movement.
V ATM = entry velocity of asteroid (or comet) to the
planetary atmosphere (m/s).
m1 , m2 = mass distributions associated to the
asteroid in movement in units of lineal density (
kg/m ), where m2 < m1.
gP = acceleration gravity on planetary surface
(m/s2).
The melt volume associated with the impact energy
for circular and elliptical craters, is determined
respectively as,
V CIRCULAR MELT (1/3) PDR(max) [ ( IMP /2)2 +
( IMP
/2 PDR(max) / tg IMP)2 + ( IMP / 2) ( IMP / 2
PDR(max) / tg IMP)] ,
and,
V ELLIPTICAL MELT (1/3) PDR(max) [ 0.894304564642
{ 12 + ( 1 2h / tg IMP)2 } + ( / 16) { 1(1 + (3 2
(2)0.5 )0.5 }{ 1 2h / tg IMP + ( 1 2h / tg IMP) (3
2(2)0.5 )0.5 } ] ,
where,
PDR(max) = maximum depth of crater.
IMP = crater diameter.
IMP = impact angle.
1 = major axis.
h = PDR(max).
These relations are deduced according the following
polynomial,
(1/32)( 2 )4 – (3/16)( 1)2( 2)
2 + (1/32)( 1 )4 > 0 ,
(1)
where,
2 = minor axis.
The expression (1) is approximately > 0 when 2 =
0.414213562 1 , i.e., for elliptical craters. When 1
= 2 > 0, i.e., for circular craters, the expression (1)
is changed to,
(1/32)( 2 )4 – (3/16)( 1)2( 2)
2 + (1/32)( 1 )4 < 0 ,
(2)
in both cases the expressions (1) and (2), take the
form of inequations.
For the calculation of macroscopical fragments [1]
(ejected) is used,
A2 H 4(NP2)4 – A2m1
4 + 2A(A – gP)m1 H 3(NP2)3 +
2A(gP – A)m13 H(NP2) = 0 ,
where,
NP2 = number of ejected fragments.
H = ( ASTEROID IN REPOSE / P) = factor of hardness.
P = planetary density.
The models used here has been applied to numerous
impact craters on both Earth and Mars, being the
results obtained very precise in the determitation of
impact conditions on the craters generated. Future
works will show more mathematical details of these
models.
References: [1] Echaurren J. C. (2007) Seventh
International Conference on Mars, Abstract # 3289.
[2] Echaurren J., and Ocampo A.C., (2003) EGS-AGU-
EUG Joint Assembly.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
GEOLOGICAL AND GEOMORPHOLOGIC STUDY OF A POTENTIAL SEDIMENTARY FAN
DEPOSIT IN ARAM CHAOS M. J. Van Helden (1,2), T.E. Zegers (2), A.P. Rossi (2), B. Foing (2), W. Van
Westrenen (1) E. Hauber (3), S. van Gasselt (4), G. Neukum (4) and the Mars Express HRSC Co-I team
Introduction: Multiple sedimentary fan deposits
have been identified in the Martian landscape.[1]Sedimentary fans are of particular interest because
of the involvement of surface fluid in their
deposition. Liquid water is an important pre-
requisite for habitability on Mars. Sedimentary fans
can either be deltas or alluvial fans; the first are
formed by sedimentation below basin water level,
the latter above water level. They can be
distinguished by different internal geometries of the
fan deposit.
Aram Chaos, located at 2.6N, 21.5W, is a chaotic
terrain in which a potential sedimentary fan was
deposited. At the eastern border, the crater is
connected to the Ares Vallis via a channel. Inprevious studies, several periods of fluid flow have
been postulated, in both directions of the channel.
At the Aram Chaos end of the Channel, a fan-
shaped unit is present, which could have been
formed by one of the fluid flow stages. The fan has
been eroded and channelled after its deposition, and
the shape of the erosional structures points to a flow
direction out of the Aram Chaos. The processes that
caused this outflow are not well understood, but
there are several suggestions such as
tectonic/volcanic activity, dome-shaping the crater
floor and declining its volume, climatic changes. In
this study the initial work of Oosthoek (2007) wasextended using HRSC images from ESA’s Mars
Express mission, and topographical MOLA data
from NASA’s Mars Global Surveyor. Data were
processed in ArcGIS and cross-sections of the
sedimentary fan deposit were made to better
constrain its geometry and formation history.
Fan formation: Several authors have suggested that
this unit is indeed a fan deposit. e.g. [1,2] Questions
we attempted to answer by geological and
morphological studies of the eastern Aram Chaos
area include whether this fan is an alluvial fan or a
delta, what its formation conditions were, what its
relationship is with other units present in the Aram
Chaos, the timing of the fan formation with respect
to other processes active in the Aram Chaos, and the
post-depositional processes that influenced the
current morphology of the fan.
The crater walls next to the fan are covered with a
relatively thin layer of sediment. This leads to the
suggestion that the formation of the fan presumably
happened in two stages: a first stage in which waterstarted to flow on top of the crater walls, depositing
a thin layer of sediments. The second stage is a
period in which the Aram Chaos channel was
incised deep enough to accommodate the amount of
water and sediment carried by the flow. In this stage
a fan-shaped deltaic structure was deposited at the
mouth of the channel.
The fan unit is eroded and channeled, most likely by
flow out of Aram Chaos. Therefore the original
geometry is only partly preserved. The advantage is
that the outflow channels provide a view into the
internal geometry of the unit. The outer morphology
of the least eroded parts of the unit was studied bycreating four cross sections.
The cross sections show the geometry of the fan.
The fan has a topographical step, which could
indicate the presence of foresets. The topographical
step is observed in all cross sections. Using HRSC
nadir images, it was possible to observe directly
indications of steep internal layering of foresets in
the erosional channels. These characteristics are
common in delta deposits. In particular, the
step/steepening of layering at the front is a typical
feature of a delta deposit as opposed to an alluvial
fan The gradient of this delta front is ~ 0.1 which is
relatively high as this slope is usually 0.01 or 0.02.This could suggest that this deltaic structure is
relatively coarse grained, as coarse-grained deltas
tend to have steeper foresets.
The cross cutting and depositional relationships
in the region have been studied in detail, in
particular using HRSC NADIR and anaglyph
images. The depositional sequence that can be
derived seems to suggest that the unit was deposited
relatively late, i.e. after crater fill and
chaotization.Post-depositional erosion has severely
modified the Aram Chaos fan. Most of the fan is
channelled, and the flow direction inferred from the
middle ground bars indicates that the most recent
flow has been out of the Aram Chaos, in the
direction of the Ares Vallis.References: [1] Cabrol and Grin (2000), LPSC XXXI ,
Abs. #1162 [2] Oosthoek et al. (2007), LPSC XXXVIII ,
Abs. #1577
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
GEOLOGIC HISTORY OF KASEI VALLES AND URANIUS DORSUM M. G. Chapman1, A. Dumke2,
E. Hauber3, G. Michael2, G. Neukum2, S. van Gasselt2, S. C. Werner4, W. Zuschneid2; 1U.S. Geological
Survey, Flagstaff, AZ, 86001, USA; 2Institute of Geosciences, Freie Universitaet Berlin, 12249 Germany;
3Institute of Planetary Research, German Aerospace Center (DLR), 12489 Berlin, Germany; 4Geological
Survey of Norway (NGU), 7491 Trondheim, Norway. [email protected]
Figure 1. East Uranius Dorsum oblique view; THEMIS VIS images over MOLA topography; view from north.
Introduction: Kasei Valles extends nearly 3000
km north from Echus Chasma (lat 1°S, long 80°W)
and turns sharply east (lat 20°N, long 75°) W. to
debouch into Chryse Planitia. Uranius Dorsum is aprominent ridge on the NW edge of Kasei that
trends NE parallel to scour marks within north
Kasei Valles. It is distinctly different in
appearance from local wrinkle ridges that trend
NW. We are mapping Kasei Valles to determine
geologic history and the origin of the channel and
unusual features like Uranius Dorsum. The
channel cuts into Hesperian material (unit Hr ) of
the Lunae Planum and Tempe Terra plateaus, and
Tharsis lava units At4 and At5 cover large parts of
the channel floor [1]. As Uranius Dorsum lies on
the floor of Kasei Valles, it postdates emplacement
of unit Hr.
Summary and Discussion: Using the
production function coefficients of Ivanov [2] and
the cratering model of Hartmann and Neukum [3]
to derive absolute ages, our crater counts of the unit
Hr indicate an average age between 3.6 to 3.8 Ga.
Ancient east-trending grooves and streamlined
islands were known to have cut Labeatis and Sacra
Mensae, remnants of high plateau material (near
25°N.) in north Kasei [1]. Our efforts show
additional east-trending streamlined islands farther
north (34°N) on the high plateau of Tempe Terra,
at least 1 km above the Kasei floor [1]. Islandheights (50-100 m) and distribution indicate the
ancient source floods from Tharsis (to the west)
were widespread and voluminous. Crater counts
indicate that this erosion took place around 2.98
Ga. On the Kasei floor, the lava plains north of
Uranius Dorsum are part of younger lava unit At4.
The At4 lavas north of the dorsum were emplaced
around 2.6 Ga. Northeast of about the mid-point of
Uranius Dorsum, unit At4 lavas show indications of
flood erosion, and here form moated areas around
Labeatis and other mensae with crater ages of 1 Ga
to 1.6 Ga [4]. South of the ridge, resistant
materials (lava flows?) that are cut by Kasei
erosion date from 1.3 Ga. The ridge area is too
small for accurate crater counts, but it likely was
emplaced between 1.3 to 2.6 Ga. Farther to thesouth, the eroded floor of Kasei is overlain by lavas
of unit At5. Emplacement of unit At 5 took
hundreds of Ma (ranging from 1.6 Ga to 90 Ma)
and overlapped relatively young episodic Kasei
floods from Echus Chasma to the south.
Uranius Dorsum is topographically much higher
than lava terminations in the area or west on the
flanks of the Tharsis rise. The Dorsum extends 2°
farther west (120 km to longitude 80°W) into the
Tharsis flank than previously mapped [1], and its
trend parallels that of ancient floods from Tharsis.
The dorsum has 58 aligned mounds along its
length; each with a central pit (Fig. 1). Some
mounds have nested pits and one shows material
extending away from its pit. The ridge has a
prominent frontal scarp on its south boundary.
Many closely-spaced incisions cut the ridge (and
frontal scarp) roughly perpendicular to its length,
and parallel to the north trend of young Kasei
Valles floods from Echus. In some places these
incisions outline streamlined blocks of ridge
material. Of the aligned ridge mounds, 47 are
breached via their central pits in a northern
direction. This trend and streamlined blocks of
ridge material suggest Uranius Dorsum predatedand was eroded by younger Kasei floods from
Echus. The frontal scarp may be due to erosion by
Kasei or ridge material abutting ice. Hypothetical
ridge origins include formation as a flood levee,
glacial moraine, littoral volcanic cones, mud
volcanoes, and fissure-fed volcanic cones. References: [1] Chapman, M. et al. (2007) LPSC
XXXVIII , Abs. #1407. [2] Ivanov, B. A. (2001) Space
Sci. Rev., 96(1), 87–104. [3] Hartmann, W. K., and G.
Neukum (2001) Space Sci. Rev., 96, 165–194. [4]
Hauber, E. et al. (2007) LPSC XXXVIII , Abs. #1666.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
GEOMORPHOLOGICAL MAPPING OF THE ELYSIUM PLANITIA REGION USING HRSC
IMAGES M. R. Balme1, J. B. Murray
1, J-P. Muller
2, J. R. Kim
2,
1Dept. of Earth Sciences, Open University,
Milton Keynes MK7 6AA, UK2Mullard Space Science Laboratory, Dept. of Space and Climate Physics,
University College London, Surrey RH5 6NT, UK. [email protected]
Introduction: We are using GIS techniques tomap the landforms of the Elysium Planitia Region
(~145-160oE, ~0-12
oN). The aim is to constrain the
processes that formed the extensive platy-ridged
terrains that characterise the region. The platy
terrain has been interpreted to have a flood lava
origin [e.g. 1, 2, 3], but other authors have suggested
a sea ice origin [e.g., 4, 5]. In this study we map the
extent of the platy terrain and determine the
topographic and stratigraphic associations with
surrounding terrain types.
We have used HRSC, MOC NA and WA,
THEMIS visible and IR, Viking, MOLA, and GRS
data during this study. The main mapping base wasa mosaic of 20 full resolution (~12.5m/pixel) HRSC
images. HRSC data are ideal for such mapping
having large areal coverage and good resolution.
Where possible, MOC NA images (1-10m/pixel)
have been used to verify contacts and surface types.
In the west of the region, Themis IR data were used
due to a lack of HRSC coverage.
The mapping reveals that the main platy terrainregion is continuous for nearly 500km and closely
follows a gravitational equipotential surface. At
lower elevations to the East are smaller regions of
platy terrain linked to the higher main region by
erosional channels. Further East the platy terrain
continues out into Amazonis Planitia. At the
Southwestern extreme of the study area, the platy
terrain terminates in a large channel system that
extends out of the study region to the West. The
platy terrain appears to be relatively young, having
very few unembayed >100m diameter craters, but it
is overlain by the Medusa Fossae Formation to the
south, suggesting a much older formation age.
References: [1] Keszthelyi, L., et al., (2000) JGR 105,
15,027-15,050 [2] Plescia, J.B. (2003) Icarus 164, 79-95
[3] Keszthelyi, L. et al., (2004) LPSC XXXV, abstract
1657 [4] Brakenridge, G.R. (1993) LPSC XXIV pp 175-
176 [5] Murray, J.B. et al. (2005) Nature 434, 352-356.
Figure 1. Example of mapping from HRSC images. Units include: Green, platy terrain and associated channels; Dark Red,lobate plains: Medium Red, smooth plains; Orange, crater materials; Yellow, Medusa Fossae Formation (hatched =
discontinuous MFF); Pink, remnant highland materials; Blue, smooth outflow channel. Background, Viking MDIM 2.1.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
A GLOBAL VIEW OVER THE MINERALOGICAL COMPOSITION OF DARK DUNES IN
MARTIAN CRATERS AND UPDATED RESULTS OF THE GRAIN SIZE ANALYSIS D. Tirsch1, R.
Jaumann1,3
, F. Poulet2, K.D. Matz
1, J-P. Bibring
2and G. Neukum
3.
1Institute of Planetary Research, German
Aerospace Center (DLR), 12489 Berlin, Germany.2
Institut d’Astrophysique Spatiale, CNRS Université Paris-
Sud, 91405 Orsay, France.3Department of Earth Sciences, Institute of Geological Sciences, Planetary Sciences
and Remote Sensing, Free University Berlin, 12240 Berlin. [email protected]
We study the Martian fine-grained dark
material by focussing on a global selection of
impact craters. These craters are interesting because
the material frequently accumulates on their floors
into huge dune fields such as barchan or transverse
dunes.
We extended our crater database up to 67
craters and updated the TES thermal inertia analysis.
The thermal inertia is a measure of a material’s
thermal response to the diurnal heating cycle and
thus expresses the resistance of a material to
temperature change. This measure is closely related
to properties such as particle size, degree of
induration, abundance of rocks and exposure of
bedrocks . Thus, we were able to reveal the grain
size of dark material dunes, also for those dune
fields of our database for that no convincing results
could be obtained so far. The dunes showing higher
thermal inertia values, which correspond to resistant
outcrops, are supposed to be consolidated [3]. The
global consideration of unconsolidated and
consolidated dunes still shows a slightly correlation
between consolidated dunes and lower elevations
(northern lowlands) and between unconsolidated
dunes and higher elevations (southern highlands).
Analysis of near infrared spectra from the
OMEGA spectrometer [4] on MarsExpress yields a
higher content of mafic unoxidized minerals such as
high and low Ca-pyroxenes and olivine. Most of the
crater deposits show strong pyroxene absorptions.
The minor part has olivine absorptions, whereas
forsterite occurs in most cases. Until this point, no
mineralogical difference between unconsolidated
and consolidated dunes could be identified. There is
also no correlation between the mineralogical
composition and the geographical location
recognisable. Pyroxene and olivine are unweathered
mafic silicates, which have never experienced a
chemical weathering. This indicates that this
unoxidized material has never had a contact to
liquid water or water vapour. Thus, mechanical
weathering could be the only process that caused the
comminution of the material [5].
In some places, a portion of a few dunes show
absorption bands of hydrated minerals indicating
that the material has underlain a alteration process.
The hydration might have been caused by the supply
of water [6], e.g. by melting H2O-frost layers.
There is no obvious correlation between
hydrated minerals and consolidated dune surfaces.
The global mineralogical distribution of the dark
dunes is shown in figure 1.
References: [1] Edgett, K.S., and P.R. Christensen
(1991), JGR 96 (E5), 22,762-22,776. [2] Mellon, M.T., etal. (2000), Icarus 148, 437-455. [3] Putzig, N.E. et al.
(2005), Icarus 173, 325-341. [4] Bibring, J.P., et al.
(2004), ESA SP 1240, 37-49. [5] Jaumann, R. (2006),
LPSC XXXVII , Abs. 1735. [6] Poulet, F. (2005), Nature
438, 623-627.
Figure 1. Global distribution of olivine and/or pyroxene-composed dark dunes
olivine pyroxene olivine and pyroxene unknown
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
GPS SYSTEM AND ITS APPLICATION IN MARS S.Roshany Yamchi1, M.Sabzeh Parvar
2.
1MSc degree in
Aerospace Engineering, Amirkabir University of Technology, Tehran, Iran.2Professor assistant, Amirkabir
University of Technology, Tehran, Iran. [email protected]
Section 1: Global positioning system is a satellite
navigation system, including a net of 24 orbitingsatellites that are in 6 orbits and in 11000 mile
distance. In fact it is a guidance and navigation
system that is made from 24 satellites. Its high
accuracy and universality is the reason that makes it
usable in various sciences. By applying GPS
system, all older systems such as: ballistic cameras,
Doppler, SECOR, LONG-C, LLR, SLR, N.N.S.S,
were gradually out of use. GPS is an operating
system that operates in all climate conditions.
Because the wave frequencies that is sent from the
GPS satellites, is at the limit of Giga Hertz and
climate condition have no effect on these waves.
The Mars future explorer whether to be automatedorbital or a human being, need a way to define their
situation. To do this important task the NASA
researchers are studying on a suitable satellite
positioning system like GPS for Mars that can also
perform as communicating network. In this article
we survey a GPS system application in Mars.
Applying GPS system in Mars will be in fact a
lunge in future robots technology for the Red Planet.
On some people’s opinion the Mars GPS systemmay seem to be luxe and unimaginable feature but
with the existence of some problems that will be
discussed in this article , using this system has a lot
of advantages and will be a lunge in the way of
complete discovery in Mars and finding out its
secrets.
References: [1] NRC (2006), Assessment of NASA’s Mars
Architecture.
[2] Lognonné et al. (2000), Planet. Space Sci. 49, 1289.
[3] Vinnik et al. (2001), GRL 28, 3031.[4] Malin et al. (2006), Science 314, 1573.
[5] Lognonné et al. (1996), Planet. Space Sci. 44, 1237.
[6] Gudkova & Zharkov (2004), Phys. Earth Planet. Int.
142, 1.
[7] Lognonné & Mosser (1993), Surv. Geophys. 14,
239.
[8] Van Hoolst et al. (2003), Icarus 161, 281.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
IS THE GYPSUM SPECTRAL SIGNATURE EXISTS IN THE JUVENTAE CHASMA ON MARS?
R.O.Kuzmin1, 2, M.V Mironenko1, N.A. Evdokimova2 . 1Vernadsky Institute of Geochemistry and Analytical
Chemistry, Russian Academy of Sciences, 19 Kosygin Str., Moscow 119991, Russia, 2 Institute for Space
Research, RAS, Moscow, 117997, Russia, [email protected]
Introduction: One of the striking OMEGA’sdiscoveries is related with detection of the gypsum
spectral signature on the upper part of the erosion
remnant (stacked of the layered deposits) within
Juventae Chasma and the spectral signature of the
kieserie and polyhydrated sulfate on the lower
flanks of the deposits [1, 2,]. However, we suppose
that the presence of the gypsum signature within the
upper suite of the layered deposits above of the
kieserite signature presence in the lower suit of the
deposits represents oneself essential contradiction
from points of view both geology and thermo-
dynamics. To clarify the question with presence of
the gypsum signature in Juventae Chasma, wereanalyzed the OMEGA spectra of Juventae Chasma
area and conducted chemical thermodynamic
modelling (based on the FREZCHEM model [3]) of
sulfates precipitation sequence at freezing and
evaporation of the hypothetical initial solution
reservoir which could exist within Juventae Chasma
in the past.
Analysis of the OMEGA data. Based on the
atmospherically corrected OMEGA spectra we
conducted mapping of the several spectral indexes
using the main adsorption bands for the gypsum
(1.75, 1.9 and 2.2 m), for the kieserite (1.6, 2.1 and
2.4 m) and polyhydrated sulfate (1.4 and 1.9 m)
on the area of the erosion remnant in Juventae
Chasma. The mapping results show that the spectral
signatures of the erosion remnant summit (as well as
of its flanks) are consistent well with the spectra of
the kieserite and polyhydrated sulfate, but not with
gypsum spectra. Moreover, in the place two main
absorption bands used for detection of the gypsum.
(1.9 and 2.2 m) are not correlated each other.
Besides, the band 1.75 m is not detectable
generally in the erosion remnant area (see Fig.1).
Results of the modelling. Because the Martian soil
enriched by such elements as Fe and Mg and
notably less by Ca [4, 5], it seems more logical that
the original solution on Mars were also enriched
mostly in Fe and Mg and less in Ca. By the reason,
the mole ratios of the main cations and anions had
been taken similar to ones in the Martian soil:
Mg/Ca=10/1 (as well as 20/1and 200/1), SO2/Cl=5,
Na/Mg=0.5, Na/K=4, Mg/Fe=5. At modelling of the
salts precipitation sequence during the evaporation
and freezing of the initial solution (at T=273K),
both the equilibrium and fractionation conditions
have been also considered. Results of the modelling
are presented on the Fig.2. As well seen from Fig.2,
at the modelled processes the gypsum beginsprecipitate first (at different ratio Mg/Ca), whereas
the magnesium and the iron sulfates and mirabilite
have been precipitated later and in much largermasses. At the equilibrium freezing of the initial
solution (from T ~263K and up to eutectic point of
solution) the deposited masses of epsomite is larger
than for gypsum ~ in 9 times and for the iron sulfate
~ in 4.5 times. The similar tendency is found at the
modelling of the equilibrium and fractional
evaporation of solution at T=273K (see Fif.2b, c). In
this way, following to modelling results the gypsum
may not to be accumulated dominantly during
formation of the last (upper) suite of the salts
deposits within the freezing (or evaporating)
solution reservoir. The sulfate could be precipitated
in the lower suites of the deposits. The smalloutcrops of the lower suites of the layered deposits
is located in several places of the Chasma’s floor [6]
and the high resolution mapping by the instrument
CRISM onboard MRO mission may to help to
identify the spectral signature of the outcrops.
Acknowledgments: This study was supported by
the Russian Foundation for Basic Research (project
N 06-0216920). References: [1]-Bibring J-P et al.,
(2005), Science 307, 1576-1581. [2] Gendrin A. et
al.,(2005), Science (Times New Roman,, 307,,1587-
1591. [3] Marion G.M et al., (2006), EOS
Trans.Agu, 87(52), Fall Meet. Suppl. [4] Wang A. et
al., (2006), JGR, 111, E2, E02S17.
Figure1. View of the erosion remnant in Juventae Chasma
(a) and the fragments of the index maps for adsorption band
1.9 m (b) and 2.2 m (c) compiled based on the OMEGA
cube ORB0482_2.
Figure2. The sulfates precipitation sequence obtained at the
thermodynamic modelling of the equilibrium freezing andevaporation (a, b) and fractional evaporation (c) of the
hypothetical initial solution.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
How to use HRSC multispectral data? Example for the Echus Chasma region. J.-Ph. Combe1, T. B.
McCord1.
1Bear Fight Center, 22 Fiddler’s Road, Winthrop, WA, 98862, USA. [email protected]
Why try to perform spectral analysis with
HRSC color data?
HRSC color data are multispectral images of Mars in the visible at 50 to 200 m/pixel. They
contain potentially high spatial resolution
information about the surface composition [1].
Visual analysis of these images is often used to
determine contrast boundaries of surface units as a
complement to morphology and texture (e.g. [2]).
The objective is to obtain spatially detailed
composition maps on the Echus Chasma region [7,
8]. Identification of surface mineralogy is performed
using near-infrared high-spectral resolution
observations such as OMEGA data at a 300 m to
several km scales. Spectral analysis of HRSC color
images is a potential way to map spectral units thatcan be related to the mineralogy at low-spatial
resolution.
Bright red and dark materials plus water ice are
the three main spectral units that can be identified in
the visible with HRSC [1, 3]. Feasibility of spectral
analysis using HRSC color data and its usefulness to
distinguish more surface types still needs to be
demonstrated. We therefore identified pre-
processing requirements, and we established a
corresponding procedure.
HRSC spectral and photometric data
HRSC multispectral data are acquired through
different wavelength filters at various angles withrespect to the nadir [4]: 440 nm (-3°), 530 nm (-3°),
650 nm (0°), 750 nm (+16°) and 970 nm (-16°).
Two stereographic and two photometric images are
also acquired at 650 nm with viewing angles of +/-
18° and +/-5° respectively. Viewing angles for
images at 650 nm are designed to calculate precise
Digital Elevation Models (DEMs) and for surface
roughness and photometric investigations (e.g. [5,
6]). The geometry of the color cameras is not for
scientific purpose, but it has to be taken into account
in spectral analysis.
Shading, shadowing and indirect illumination
The different viewing angles must be considered
when interpreting the data. Topography at scales
smaller than a resolution cell (or surface roughness)
creates surface shade and projected shadows.
Proportion of shaded areas depends on geometrical
characteristics of the surface that can be described
by statistics of slopes, heights and their spatial
distribution [9]. It varies also with the geometry of
illumination and observation such as incidence
angle, difference between azimuth of observation
and solar azimuth and emergence angle [10].
Atmospheric scattering may contribute in
different ways, depending on the phase angle. [11]
found that aerosol scattering can contributes to 25%
of the signal for zero-phase angles at 1 m on dark
areas, even with a clear atmosphere. This amount
reaches a maximum at low phase angle but is still
significant for lower angles, and it varies withwavelength in a non-linear way.
As a consequence, the amount of shade in a
given pixel varies from image to image in HRSC
color data, and this affects the shape of spectra. The
five HRSC panchromatic data can be used to derive
photometric information. However, the number of
unknown parameters exceeds the number of bands.
Evaluation of the most important parameters
According to [11], contribution by aerosol
scattering is significant enough to be considered.
We used simulations to investigate photometric
properties related to the surface at a HRSC pixel
size scale. Measured DEMs are used as templates,and they can be also modified numerically to
simulate different types of surfaces as well as
geometries of illumination and observation.
Projected shadows are calculated. Models of
scattering from surrounding assume each element of
the surface is Lambertian, like in [12]. Comparisons
are made specifically according to the HRSC
viewing angles. A spectrum of a projected shadow
is used as an empirical correction for aerosols
scattering.
First results: Does Mars appear simply red?
Incidence angle, emergence angle, and slopes of
a rough surface are the main parameters that affectHRSC color data. Signal variations due to projected
shadows may be above the noise level when
incidence angles are high. Difference in azimuth of
illumination and observation is significant for very
high incidence angles and very rough surface only.
We performed Spectral Mixture Analysis (SMA)
on HRSC images after roughness corrections.
Spectral endmembers are bright red and very dark
components, plus the aerosol scattering
contribution. SMA residuals on the Echus Chasma
region are below the noise level for all the pixels,
meaning two spectral surface components are
sufficient to model all units with five bands in the
visible. The resulting spectral map will be shown
and interpreted. Further investigations are currently
in the process in order to refine these results and
provide quantifications.References: [1] McCord, T. B. et al. (2007), JGR 112,
doi:10.1029/2006JE002769. [2] Loizeau, D. et al. (2007)
JGR 112, doi:10.1029/2006JE002877. [3] Hauber E. and
Neukum G. (2006), Astronomy Geophysics 47. [4]
Neukum, G. et al. (2004), ESA SP , 1240. [5] Mushkin et
al. (2006), GRL 33, doi:10.1029/2006GL027095. [6] Cord
et al. (2007), Icarus, in press. [7] Masson et al. (2005),
LPSC no. 1340. [8] Chapman et al. (2007), LPSC no.
1338. [9] Shepard et al., 2001). [10] Combe et al. (2007),
LPSC no. 2367. [11] Erard (2001), GRL 28. [12]
Schkuratov et al. (2005), Icarus 173.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
HRSCview: ONLINE ACCESS TO MARS EXPRESS HRSC IMAGE AND DTM MOSAICS. G.
Michael, A. Dumke, S. Walter, G. Neukum, Freie Universitaet Berlin, Germany. (gregory.michael-at-fu-
berlin.de)
The High Resolution Stereo Camera (HRSC) on
the ESA Mars Express spacecraft has been orbitingMars since January 2004. By spring 2007 it had
returned around 2 terabytes of image data, covering
around 35% of the Martian surface in stereo and
colour at a resolution of 10-20 m/pixel. HRSCview
provides a rapid means to explore these images up
to their full resolution with the data-subsetting, sub-
sampling, stretching and compositing being carried
out on-the-fly by the image server. It is a joint
website of the Free University of Berlin and the
German Aerospace Center (DLR).
The system operates by on-the-fly processing of the six HRSC level-4 image products: the map-
projected ortho-rectified nadir pan-chromatic and
four colour channels, and the stereo-derived DTM
(digital terrain model). The user generates a request
via the web-page for an image with several
parameters: the centre of the view in surface
coordinates, the image resolution in metres/pixel,
the image dimensions, and one of several colour
modes. If there is HRSC coverage at the given
location, the necessary segments are extracted from
the full orbit images, resampled to the required
resolution, and composited according to the user’s
choice. In all modes the nadir channel, which has
the highest resolution, is included in the composite
so that the maximum detail is always retained. The
images are stretched according to the current view:
this applies to the elevation colour scale, as well as
the nadir brightness and the colour channels. There
are modes for raw colour, stretched colour,
enhanced colour (exaggerated colour differences),
and a synthetic ‘Mars-like’ colour stretch. A colour
ratio mode is given as an alternative way to examine
colour differences (R=IR/R, G=R/G and B=G/B).
The final image is packaged as a JPEG file andreturned to the user over the web. Each request
requires approximately 1 second to process.
A link is provided from each view to a data
product page, where header items describing the full
map-projected science data product are displayed,
and a direct link to the archived data products on the
ESA Planetary Science Archive (PSA) is provided.
At present the majority of the elevation composites
are derived from the HRSC Preliminary 200m
DTMs generated at the German Aerospace Center
(DLR), which are not available as separatelydownloadable data products. These DTMs are being
progressively superseded by systematically
generated higher resolution archival DTMs [1], also
from DLR, which will become available for
download through the PSA, and be similarly
accessible via HRSCview. At the time of writing
this abstract (September 2007), 155 such high
resolution DTMs are available for download via the
HRSCview data product pages.
We are beginning work on producing HRSC
image and DTM mosaics for selected regions of the
Martian surface. It is planned that HRSCview will
additionally provide access both for browsing and
downloading these new products.
References [1] K. Gwinner, F. Scholten, R. Jaumann,
T. Roatsch, J. Oberst, G. Neukum. Global mapping of
Mars by systematic derivation of Mars Express HRSC
high-resolution digital elevation models and orthoimages
ISPRS IV/7 Extraterrestrial Mapping Workshop, Houston,
2007.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
AN ICE DYNAMICS ORIGIN FOR MARTIAN ICE CAP GEOMORPHOLOGY: NEW DATA FOR A
VIKING ERA INVESTIGATION. J.H.J Leach1.
1University of Melbourne, Parkville, Vic. 3010, Australia..
Section 1: In 1979 a project was undertaken at
the University of Melbourne to model the NorthPolar Ice Cap of Mars using techniques developed
for the ice sheets of Greenland and Antarctica (Budd
et al. 1986). This model was based on the
assumption of “steady state”: that the accumulation
on the central part of the ice cap was matched by the
mass lost by ablation from the “layered terrain”.
This assumption allows the estimates of the mass
gain and loss across the ice cap and the age of the
ice could be calculated by integration of velocities
along the flow lines. The model gave a maximum
ice movement of 50cm horizontally and 20cm
vertically, although the average rates were lower
than this by two orders of magnitude. The totalinferred mass gain and loss was of the order of
0.2km2
per Mars year which would represent about
20% of the atmospheric water turnover calculated
by Jakosky and Farmer (1982). The calculated
residence time of the ice was up to 100 million
years.
The model assumed that the peripheral “layered
terrain” was in fact an ablation zone of the ice cap
itself and not some underlying sedimentary deposit.
using this assumption, it was possible to explain all
of the observed geomorphic features of the ice cap
in a way which was consistent with the output of thenumerical model (Leach 1982a,b 1983a,b). Newer,
high resolution imagery suggests that this
interpretation of the geomorphology is in fact
correct.
However, the model could only really calculate
in orders of magnitude since the quality of the
available data imposed real restrictions on accuracy
and precision. Key among these limitations were the
lack of detailed topographic information across the
ice cap and that the model assumed a thickness of
ice obtained by extrapolation from the surrounding
terrain. In both cases newer data could be used to
update this model so as to provide a more precisemap of ice movement. This in turn could be used to
test the correlations between ice movement and
polar geomorphology which formed the basis of the
earlier interpretations of geomorphic process.
Updating this work is important since the
question of the nature and origin of the polar
laminated terrains is crucial to considerations of the
timing and scale of any past changes in the Martian
climate.
References: Budd, W.F., Jenssen, D., Leach, J.H.J.,
Smith, I.N., and U. Radok 1986. The North Polar Ice Capof Mars as a Steady-State System. Polarforschung 56
(1/2)
Jakosky, B.M. and C.B. Farmer 1982. The seasonal and
global behavior of water vapor in the Mars atmosphere:
Complete global results of the Viking Atmospheric Water
Detector Experiment. J. Geophys. Res. 87(84): 2999-
3019.
Leach, J.H.J. 1982. "An Ice Dynamics Origin for the
Martian Polar Laminated Terrains." Lunar and Planetary
Science Conference XIII, Lunar and Planetary Institute,
Houston.
Leach, J.H.J. 1982. "The Development of Elongate
Structures within the Martian Polar Ice Cap." Lunar and
Planetary Science Conference XIII, Lunar and Planetary
Institute, Houston.
Leach J.H.J. 1983. "The Rapid Ablation Zone - A
Distinctive Terrain unit in the Martian North Polar Ice
Cap." Lunar and Planetary Science Conference XIV,
Lunar and Planetary Institute, Houston.
Leach J.H.J. 1983. "Stagnant Ice in the Martian North
Polar Ice Cap." Lunar and Planetary Science Conference
XIV, Lunar and Planetary Institute, Houston.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THE IMPACT OF DEGASSING ON THE EVOLUTION OF THE ATMOSPHERE OF MARS. C.
Gillmann1, P. Lognonné
1.
1 Institut de physique du Globe de Paris, 4 avenue de Neptune 94100 Saint Maur,
France. [email protected]
We study possible states of the past Martian
atmosphere consistent with present observationthrough a simple evolution model based on realistic
outgassing scenarios.
Degassing is a direct consequence of the amount of
melted material created by the activity of the mantle
of the planet. Here, crust production rates from
models such as that by Breuer et al. (2006) or
Manga et al. (2006) are taken as input for the mantle
degassing. The evolution of the volatile contents of
the atmosphere is studied through different
scenarios. Since the other effect that has a strong
influence on the atmosphere is its escape into space,
it is also taken into account. Hydrodynamic escape
mostly takes place during the first few hundreds of million years so other processes for atmospheric
escape have been considered in order to quantify the
loss of volatiles during later periods. In the case of
Mars in particular, where no evidence for carbonates
has been found, escape seems to be the main
mechanism for CO2 removal. Using data from Mars
Express and several models such as created by
Leblanc (2001) and Chassefière, Leblanc and
Langlais (2006), a model for the evolution of
Martian atmosphere and volatiles has been set up.
We first focused on the present situation as
described by available data such as those from Mars
Express in order to study the late evolution of theMartian atmosphere. It appears that a crustal
production of at least 0.01 to 0.1 km3 /year is needed
for the atmosphere to be at steady state at present-
times which is consistent with low activity as
observed today.
Our study focussed mainly on the evolution of CO2,
and results for several scenarios show that when
degassing is intense enough (which is likely given
the range of mantle compositions we investigated),
most of the present atmosphere of Mars would be of volcanic origin rather than some residue from a
primordial one that would have been depleted by
atmospheric escape over geological times. This
means that the present-day atmosphere is rather
recent. Our models strongly imply that in most of
the cases the atmosphere of Mars is not much older
than 1.5 Gyr, with some cases where it is as young
as 1 Gyr. This is true for both our degassing data
sets and with a wide range of parameters.
Our study also provides us with rough constraints on
the CO2 concentration in the lavas of Mars, showing
that relatively low concentrations are compatible
with the present situation but that Earth-likequantities would probably be incompatible with the
atmosphere we observe today.
Finally, such modelling can give insight on the
state of the atmosphere of Mars in the past, which
might be instrumental in explaining some features
that are observed on the surface of the planet such as
sulphates and phyllosilicates detected by OMEGA
(Bibring et al., 2005) and that have very special
formation conditions.
References: Bibring, J.-P., et al. (2005). Science 307,
1576-1581.
Breuer D., Spohn T., 2006. Planetary and Space Science 54 (2006) 153–169.
Chassefière, E., F. Leblanc and B. Langlais (2006),
Planetary and Space Science Volume 55, Issue 3, Pages
343-357
Leblanc F., and R.E. Johnson, (2001), Planetary and
Space Science, Volume 49, 645-656.
Manga, M. et al., (2006)., AGU fall meeting, Abstract
#P31C-0149.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
IONOPAUSE FEATURES OF MARS AS OBSERVED BY THE RADIO SCIENCE EXPERIMENT
MARS ON MARS EXPRESS K. Peter (1), M. Pätzold (1), B. Häusler (2), S. Tellmann (1) and G.L. Tyler
(3)(1) Rheinisches Institut für Umweltforschung, Abteilung Planetenforschung, Universität zu Köln, Cologne,
Germany, (2) Institut für Raumfahrttechnik, Universität der Bundeswehr München, Neubiberg, Germany, (3)
Department of Electrical Engineering, Stanford University, Stanford, California, USA
The ionopause of a planet is defined as the boundary
between the ionosphere and the solar wind regime.
It was first described for Venus when a sharp
decrease in electron density towards very small
values was found at certain altitudes. So far, the
ionopause at Mars has not been well observed. One
reason is that the noise of the Viking profiles was
relatively high and did not drop below 500 el/cc.
The MGS data base is inconclusive concerning the
ionopause.
The highly elliptical orbit of Mars Express
allows us to investigate the electron density of Marsup to an altitude of about 1500 km. We want to
define the ionopause feature at Mars as an electron
density gradient starting well above the topside
ionospheric main peak, tending to decrease the
electron density towards noisy values around zero.
The Radio Science Experiment MaRS on
Mars Express sounded the Martian atmosphere and
ionosphere during four occultation seasons starting
from April 2004. So far, more than 400 vertical
profiles of the ionospheric electron density could be
derived covering both hemispheres and almost all
local times. This presentation will show the high
variability of the ionopause structures of Mars.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
LARGE BASINS OF WATER ON MARS V.V. Yakovlev Kharkiv national academy of municipal economy
61001, 38, Kirova str., Kharkiv, Ukraine. [email protected]
On the surface of Mars, previously known
structures were decoded like large ice formations of
hydrolaccoliths that contain sizeable stores of iceand liquid water. In the paper author contribute
cases of ice essence of this structures and existence
of liquid water inside structures.
Powerful water head in hydrolaccoliths evince
thickness of permafrost body under the
hydrolaccoliths at issue.
The formation of hydrolaccoliths suggests
existence of artesian water resources.
The tops of these structures consist, probably, of desalinated ice. For the principal scheme of the
structure of ice storages of water see fig. 1.
Under the ice surface there exist liquid water
reservoirs which periodically erupt though,
reaching the surface. The water, in all probability is
mineralized. It probably contains silicic acid and
iron.
• Figure 1. Scheme of a possible structure of a Martian hydrolaccolith. 1. The crystal base. 2. Sedimentary rocks.3.
Mineralized ice. 4. Fresh ice. 5. Spots of outpouring of the water. 6. Boundary of permafrost. 7. A head surface of water. 8.
Water reservoir. 9. The gullies. 10. A direction of movement underpermafrost waters. 11.A direction of water flow on the
surface of hydrolaccolith.
After the discovery of Martian hydrolaccoliths
there appeared a need to refine and revise the
nature of many structures and hillside processes on
the surface of Mars, i.e. the causes of the former
Martian floods. In the paper, author polemize with
opponents, which are the experts of NASA, on the
question of essence dark and light slope streaks.
In case of destruction of hydrolaccoliths significant
amounts of water both from internal storages, and
from resilient and capacious storages of artesian
waters will pour out on the surface. Probably, such
breaks also caused in the past the catastrophic
outflows of the Martian rivers. The very existence
of the cryosphere on Mars triggers catastrophic
floods, as the probability of earthquakes and falling
of meteorites always exists.
Both desalinated and mineral ice of hydrolaccoliths
has a great value for the future Martian generations
of settlers as a source of water, mineral resources
and heat. However on account of danger of
catastrophic destruction and flooding, which might
be caused by careless disrupt of an equilibrium
system, these objects should be subjected to
profound examination. Author proposes make these
objects as high-priority for the program of Mars
studding by unmanned device.
In our opinion, the examination of liquid water of
hydrolaccoliths with a view of forms of life might
have greatest scientific value. Up until now the
most promising object of this kind was Europe -
the satellite of the Jupiter. But it appears the from
now on Mars will take its place.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Laser Based Dust and Wind Flow Sensor: DDES Part of the MEDUSA Dust Suite.
J. P. Merrison1, L. Colangeli
2, J. J. Lopez-Moreno
3, P. Nørnberg
1, P. Palumbo
4, F. Esposito
2, J. F. Rodriguez-
Gomez3.
1Mars Simulation Laboratory, Aarhus University, Ny Munkegade bygn 1520, Aarhus 8000C,
Denmark,2INAF-Osservatorio Astronomico di Capodimonte, Via Moiariello 16 – 80131 Napoli, Italy.
3Instituto
de Astrofisica de Andalucia, Camino Bajo de Huetor 24 – 19008 Granada, Spain, 4Università degli Studi diNapoli "Parthenope", Dipartimento di Scienze Applicate, Via A. De Gasperi 5 – 80133 Napoli, Italy
Abstract
Wind borne dust is the most active environmental
factor affecting the Martian surface and its
atmosphere, yet there still lacks a detailed physical
understanding how it is transported.
The DDES is a miniature, laser based instrument
which integrates sensors capable of quantifying
important parameters needed for the understanding
and modeling of dust transport on Mars, these
include: wind speed, wind direction, suspended dustconcentration, dust deposition and removal rates as
well as the electrification of the Martian dust [1,2].
As part of the MEDUSA dust suite it gives ExoMars
the most thorough dust analysis package to land on
Mars.
Laser based wind sensing has the great advantage of
being non-contact and allows flow to be quantified
independent of the pressure, temperature,
composition and so on. It gives the capability, for
the first time, to directly quantify the suspended dust
concentration close to the Martian surface.
Of importance both scientifically and
technologically is the deposition rate of dust, this ismeasured by the DDES using light scattering from a
surface. By applying electric fields dust
accumulation can be enhanced and the
electrification of suspended dust grains quantified.
Electrification plays an important role in the physics
of granular material, specifically it can be involved
in adhesion and cohesion i.e. aggregate formation
[3,4]. Although dust electrification has been seen
from experimental simulations to be of considerable
importance to dust transport on Mars, this would be
the first measurement of the process on Mars.
Figure 1. On the left a photograph shows a view inside
the Aarhus Mars Simulation wind tunnel while exposing
the DDES prototype sensor to suspended Martian dust
analog. On the right a photograph taken after the exposure
showing enhanced dust accumulation on the transparent
electrodes.
The combined functionality of the DDES instrument
potentially allows the diurnal and seasonal dust
cycle to be quantified and understood on a physical
level. The instrument is low mass (<100g), low
power (<1W) and robust. This is made possible by
utilizing recent advances in optoelectronics. Testing
of a prototype instrument has been performed undersimulated Martian conditions in a wind tunnel
facility [2,5].
There is still much to be learned about granular
transport and especially the role of grain
electrification on earth. In this respect this
instrument is also used routinely as a laboratory
(and field test) sensor for terrestrial granular
transport studies.
References:
1. ”A Miniature Laser Anemometer for Measurement
of wind speed and dust suspension on Mars”,
Merrison, J.P.; Gunnlaugsson, H.P.; Jensen, J.;
Kinch, K.; Nørnberg, P.; Rasmussen, K.R. (2004)
Planetary and Space Science; 52(13): 1177-1186
2. “An integrated laser anemometer and dust
accumulator for studying wind induced transport on
Mars”, Merrison, J.P., Gunnlaugsson, H.P., Kinch,
K., Jacobsen, T.L., Jensen, A.E., Nørnberg, P.,
Wahlgreen, H. (2006) Planetary and Space Science
(2006), 54, 1065-1072
3. “The electrical properties of Mars analogue dust.”,
Merrison, J.; Jensen, J.; Kinch, K.; Mugford, R.;
Nørnberg, P. (2004) Planetary and Space Science;
52: 279-290
4. “Determination of the Wind Induced Detachment
Threshold for Granular Material on Mars using Wind
Tunnel Simulations.”, Merrison, J.P., Gunnlaugsson,H.P., Nørnberg, P., Jensen, A.E., Rasmussen, K.R.,
Icarus, accepted 2007
5. ”Simulation of the Martian Aerosol at Low Wind
Speeds”, Merrison, J.P. et al. (2002) JGR, 107, 16-1
to 16-8, 2002
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Laser induced breakdown spectroscopy of soils and rocks under Martian conditions
I. Rauschenbach1, S. G. Pavlov
2, V. Lazic
3, H. W. Hübers
2and E. K. Jessberger
1.
1Institut für Planetologie,
Wilhelm-Klemm-Str. 10, 48149 Münster, Germany.2DLR, Institut für Planetenforschung, 12489 Berlin-
Adlershof, Germany.3
ENEA, FIS-LAS, Via Enrico Fermi 45, Frascati, Italy. [email protected]
Introduction: ExoMars, ESA´s upcomingmission to Mars will include a combined
Raman/LIBS instrument. The LIBS part will
analyze the elemental composition of Martian
surface rocks and soils in the temperature range
from +30°C to -60°C [1]. Results of various Mars
observations infer that sorption water is a soil
constituent in the upper meters of Martian surfaces
at mid- and low latitudes [2,3]. Consequently, the
presence of pore and adsorption water including
their transformation into ice phases and vice-versa
must be considered in the LIBS analyses of Martian
soils and rocks. In the literature we find only few
LIBS surveys on ice [4] and on soil/ice mixtures [5].Therefore, we have started systematic LIBS
analyses of wet samples at variable surface
temperatures [6,7]. The present work is based on the
GENTNER project [1] and is funded by German
Aerospace Center DLR. Here we report new results
of our studies for the Raman/LIBS instrument for
ExoMars on the dependence of LIBS signals from
relevant sample types under Martian environmental
conditions as a function of sample surface
temperature.
Experimental: We used a Nd:YAG laser
operated at 1064 nm with 8 ns pulse width and
10 Hz repetition rate. The spectrometer for theplasma emission is a high-resolution Echelle mono-
chromator equipped with a gated ICCD. The sample
holder can be cooled with liquid nitrogen to -80°C
in pre-selected steps. The chamber is filled with a 7
mbar “Martian” atmosphere (95.55% CO2, 2.7% N2,
1.6% Ar, 0.15% O2).
Results: We analyzed andesite rock samples
featuring different grades of surface roughness/pore
sizes and pressed powder pellets of the same
samples and of certified reference materials.
The LIBS signal from all samples shows strong
drops below 0°C. We attribute this characteristic
signal behavior to the presence and to phase
transitions of (supercooled) water on surface grains
and inside surface pores and scratches [6,7]. The
specific transition temperatures depend on surface
roughness and pore size. Three main transition
temperatures of water, supercooled water and water
ice had previously been established around 0°C, -
40°C and -50°C [8-13]. They are also observed with
the LIBS technique as sharp signal dips (Fig. 1a). At
0°C water inside larger pores and scratches
nucleates to hexagonal ice. When this free water is
slowly cooled down it can exists as supercooled
water down to -40°C. At this temperature the
supercooled water nucleates homogenously to cubic
ice. On areas with gently corrugated surface
supercooled water can exist down to -80°C. Around-50°C it changes its thermodynamic properties; a
transition from normal liquid structure to an
amorphous hydrogen-bonded network is
hypothesized [12]. As shown in Fig. 1b, the Si/H
minima are directly correlated with increased
hydrogen emission peak intensities. This result
strongly corroborates the hypothesis.
Conclusions: LIBS signals from water bearing
Martian analogue samples under Martian conditions
in the range +30°C to -60°C are a function of
sample temperature. We observed signal drops
below 0°C and attribute this behavior to phase
transitions of supercooled water present inside thesurface pores and scratches. On one hand, this effect
might significantly influence the analytical
capability of the LIBS technique. But on the other
hand it might allow measuring the water content on
soil and rock surfaces and inside their pores.
-60 -50 -40 -30 -20 -10 0 10 20
0.2
0.4
0.6
0.8
1.0
n o r m a l i z e d i n t e n s i t y
Temperature (°C)
HSi/H
-60 -50 -40 -30 -20 -10 0 10 200.8
0.9
1.0
a)
S i i n t e n s i t y
b)
heating cycle
30 J/cm2
Figure 1. a) Normalized LIBS Si peak intensity (288.2
nm) on smooth andesite rock sample and b) Normalized H
peak intensity (656.2 nm) and Si/H ratio as a function of temperature.
References: [1] Jessberger et al. (2003) ESA Call for
Ideas of the Pasteur instrument payload for ExoMars
rover mission. [2] Mitrofanov et al. (2003), LPSC 34,
1104. [3] Möhlmann, D. (2005), LPSC 36 , 1120. [4]
Càceres et al. (2001) Spectrochim. Acta B, 56, 831-838.
[5] Arp et al. (2004) Appl. Spectrosc., 58, 897-909. [6]
Rauschenbach et al. (2007) LPSC 34, 1284. [7] Lazic et
al., Spectrochim. Acta B (submitted). [8] Dash et al.
(1995) Rep. Prog. Phys., 58, 115-167. [9] Engemann et al.
(2004) Phys. Rev. Lett., 92, 205701:1-4. [10] Schreiber et
al. (2001) Phys. Chem. Chem. Phys., 3, 1185-1195. [11]
Bergman et al. (2000) J. of Chem. Phys., 113, 357-363.
[12] Debenedetti et al. (2003) Phyics Today, 56, 6 , 40-46. [13] Ito et al. (1999) Nature, 398, 492-495.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
LAYERED DEPOSITS OF THE EASTERN VALLES MARINERIS AND CHAOTIC
TERRAINS ON MARS M. Sowe1, E. Hauber
1, R. Jaumann
1, 2, K. Gwinner1, F. Fueten3, R. Stesky4, and G.
Neukum2 1Institute of Planetary Research, German Aerospace Center (DLR), Berlin, Germany,
2Department of
Earth Sciences, Institute of Geological Sciences, Planetary Sciences and Remote Sensing, Free University
Berlin, Berlin, Germany,3Department of Earth Sciences, Brock University, St. Catharines, Ontario, Canada,
4Pangaea Scientific, Brockville, Ontario, Canada. [email protected]
Light-toned and layered deposits (LDs)
are present throughout the whole Valles Marineris
and adjacent chaotic terrains. They are supposed to
be of sedimentary [1,2,3] or volcanic origin [4].
Using high-resolution image and elevation data, we
study their morphology, elevation, thickness, layer
geometry, and consolidation in order to ascertain
how they formed.
LDs show differing morphologies. There
are light-toned mounds with a flat top and steep
slopes, flow-like structures where light-toned
material flows around the chaotic-terrain material
as well as terrace-like structures and razorblade-
shaped morphologies that show massive cap rocks
at their top and layering in lower parts. Often, there
is a diffuse contact between LDs and chaotic terrain
due to dust coverage. Many of the LDs exhibit
yardangs, suggesting weakly consolidated and fine-
grained material shaped by wind erosion. Wind
activity is also indicated by dunes that occur in
depressions (e.g., fractures) that are cut into the LD
surfaces. Debris fans and a general lack of boulders
at their base may indicate loose to partly
consolidated sedimentary material. This is alsoconfirmed by TES-derived thermal-inertia values
of ~ 300 SI indicating rock materials [7].
Elevation data show that LDs are located
in depressions at different elevations but far
beneath the surrounding plateau rims (1000-4000 m
in the chasmata, 200-1500 m in the chaotic
terrains). LDs are superimposed on chaotic-terrain
material and are therefore younger. Strike and dip
measurements point towards sub-horizontal
layering (in the range of < 10°) and NS- to NNE-
SSW-strike for Iani Chaos. LD thicknesses vary in
the range of 200-4000 m, assuming ILDs have
horizontal to sub-horizontal stratification (Fig. 1.1,
1.2).
When looking at higher-resolution MOC
images, deposits show varying surfaces (rough,
fractured, grooved, cap rock). Different surfaces
textures may be due to differences in consolidation
and/or wind erosion; the mineralogical composition
is however comparable. LDs are closely connected
to sulphate- [5] and hematite rich materials [6]. A
topographic trend is observed as some LDs show
surfaces that are restricted to chaotic terrains and
other to chasmata.
References:
[1] Peterson, C. (1981), Proc. Lunar Planet. Sci. Conf.,
11th, 1459-1471. [2] Nedell et al. (1987), Icarus, 70,
409-441. [3] Malin M. C., and K. S. Edgett (2000),
Science, 290, 1927-1937. [4] Chapman M. G. (2002),
Geol. Soc. Spec. Publ., 202, 273-303. [5] Gendrin A. etal. (2005), Science, 307 , 587-1591. [6] Glotch T.D., and
P.R. Christensen (2005), JGR, 110,
doi:10.1029/2004JE002389. [7] Putzig et al. (2005),
Icarus, 173, 325-341.
Fig. 1.1: MOLA-map showing the locations of LDs in the research area (red circles).
Fig. 1.2: LD thicknesses from west to east (sub-horizontal layering assumed).
0
500
1000
1500
2000
2500
3000
3500
4000
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
LAYERED MORPHOLOGY OF THE LATITUDE-DEPENDENT MANTLE. S. C. Schon1, J. W. Head
1
and R. E. Milliken2
1Dept. of Geological Sciences, Brown University, Providence, RI, 02912 USA.
2Jet
Propulsion Lab, 4800 Oak Grove Dr., Pasadena, CA 91001 USA. [email protected].
Introduction: Systematic latitudinal variations
in surface roughness associated with ice-richmantling deposits have been documented using
MOLA [1]. Using MOC images, [2,3,4] presented
morphological observations of young surface
textures ranging from smooth and continuous to
highly degraded that they interpreted as consistent
with the recent emplacement of ice-cemented loess
undergoing desiccation/degradation. Modeling
results [5,6]) and gamma-ray spectroscopy [7,8])
support the stability of near-surface ground ice in
this latitude regime. Stratigraphic analysis of
layering within mantle deposits serves as a means
of assessing formation hypotheses such as vapor
diffusion (e.g., [9]) as well as airfall deposition thatmay be correlative with geologically recent
obliquity perturbations; Figure 1.
Figure 1: A) Strict vapor diffusion (left) and obliquity-
driven surface emplacement of mantling materials
(right). B) High obliquity expands the ice stability zone.
Setting: Layering within the mantle is
observed symmetrically within the mid-latitudes of
both hemispheres, especially in the southern
hemisphere where topographic variability more
frequently generates favorable slopes that exposelayering. Mantle surface texture varies from
smooth and continuous at higher latitudes to
discontinuous degraded mantle textures at lower
mid-latitudes. Layering outcrops are concentrated
in the transitional zone between these textures
(~35°-40°); both smooth and degraded mantle
textures are observed in MOC (101) and HiRISE
(42) images that contain mantle layering outcrops
(x38°S). This latitudinal range is commensurate
with a band of strong slope asymmetry attributed to
obliquity-controlled insolation geometry that
favored downslope movement on pole-facing
slopes [10]. Smooth mantle textures are observed
preferentially on equator-facing slopes, while
degraded mantle textures exhibit a preference for
pole-facing slopes; asymmetrically mantled craterswhich illustrate this phenomenon are common in
the transitional zone between textures.
Layer Morphology: Individual layers are
interpreted to be of relatively uniform thickness (on
the order of several to ten meters). Cross-bedding
relationships are not observed; therefore, layers are
interpreted to be of wide aerial extent.
Approximately tripartite layering with degraded
facies is the most common style of layering outcrop
observed; however, finer layering is observed on
more gentle slopes where 8 individual layers are
sometimes distinguishable; Figure 2. These units
are discernable by slight variations in texture andalbedo near the limit of resolution in MOC data.
Within the sequence, smooth textured lower albedo
surfaces separate higher albedo, raised relief,
blocky and segmented surfaces. These blocky
segments are interpreted as semi-consolidated dust-
rich lag deposits.
Figure 2: MOC M0204329 (38.01°S, 113.59°W): fine
mantle layering on a gentle pole-facing slope.
Implications: Layered mantle outcrops suggest
syndepositional layering that may be the result of
cyclical deposition of an ice-rich eolian dust
material. This is inconsistent with a strictly vapor
diffusion model of high latitude terrain softening.
Individual units are hypothesized to represent
geologically recent obliquity excursions that
mobilized volatiles for mid-latitude deposition.References: [1] Kreslavsky, M., and J. Head (2000), JGR
105, doi:10.1029/2000JE001259. [2] Mustard, J. et al.
(2001), Nature 412, doi:10.1038/35086515. [3] Milliken, R.
et al. (2003), JGR 108, doi:10.1029/2002JE002005. [4]
Milliken, R. and J. Mustard (2003), 6 th
Int. Conf. on Mars,
Abs. #3240. [5] Mellon, M. and B. Jakosky (1995), GJR
100, doi:10.1029/95JE01027. [6] Bandfield, J. (2007),
Nature 447, doi:10.1038/nature05781. [7] Boynton, W. et al.
(2002), Science 297, doi:10.1126/science.1073722. [8]
Feldman, W. et al. (2004), JGR 109, doi:10.1029/2003JE002160. [9] Chamberlain, M. and W.
Boynton (2007), JGR 112, doi:10.1029/2006JE002801. [10]Kreslavsky, M. and J. Head (2003), GRL 30,
doi:10.1029/2003GL017795.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
LMD-MGCM extended up to the thermosphere: capability for the study of Mars Express upper
atmosphere data. F. González-Galindo1, F. Forget1, M.A. López-Valverde2
, M. Angelats i Coll1
.
1
Laboratoire de Météorologie Dynamique, Université Pierre et Marie Curie, 4 Place Jussieu, Paris, France.2Instituto de Astrofísica de Andalucía, CSIC, Camino Bajo de Huétor, Granada, Spain, [email protected]
Model description The General Circulation
Model for the Martian atmosphere developed at the
Laboratoire de Météorologie Dynamique
(LMD-MGCM) in the 90s [1] has been recently
extended up to the thermosphere in collaboration
with the Instituto de Astrofísica de Andalucía
(Spain), in the frame of a project sponsored by ESA
and CNES. For this extension, new
parameterizations for physical processes appropriate
for thermospheric altitudes (in particular the NLTE
corrections to the IR radiative transfer by CO2, the
absorption of UV solar radiation and a chemicaltime-marching code, the thermal conduction and the
molecular difusion [2, 3, 4]) have been included in
the model. In this way, the LMD-MGCM has
become the first GCM for Mars able to study in a
self-consistent way the whole altitude range from
the surface to the thermosphere. This is a very
important fact for the study of the upper Martian
atmosphere, given the strong coupling between the
lower and the upper atmosphere that some recent
data have shown (e.g. density measurements during
MGS aerobraking [5], thermospheric polar warming
detected by Mars Oddyssey [6], NO nightglow
observed for the first time by SPICAM [7]). Ourextended model is able to capture these couplings
between different atmospheric layers and between
different processes. The LMD-MGCM is, thus, a
powerful tool for the study of the upper Martian
atmosphere and for the analysis of data from Mars
Express regarding this region (e.g. SPICAM and
OMEGA data), as well as from future missions.
Validation activities. The results of our
thermospheric GCM have been subject to a number
of validation activities, which include extensive
internal consistency tests and comparison with other
model and with a few satellite measurements. In
particular, an intercomparison campaign with the
Mars Thermospheric GCM (MTGCM, [8]) has been
made. During this campaign the thermal and wind
structure of the upper atmosphere as well as its
variability with season and with the dust amount in
the lower atmosphere have been tested in both
models, using the same basic input parameters. A
good general agreement is found, although we have
identified some differences at small scales [9]. Applications We are using the model to analyze
the most recent satellite observations of the upper
Martian atmosphere. In particular, we are comparing
the thermal profiles given by the model with the
temperature measurements by SPICAM using stellar
ocultation [10]. A systematic overestimation of the
lower thermospheric temperature is found, as well
as an underestimation of the altitude of the
mesopause. We are currently investigating the
thermal balance at mesospheric altitudes to find out
clues to this behaviour.
We have also studied the thermospheric polar
warming detected by Mars Odyssey during Northern
winter [6]. Our model reproduces a warming under
such conditions and sensitivity tests that show that
the in-situ tides play an important role in this polar
warming. The variability with the dust amount in
the lower atmosphere has also been studied.We also plan to use the LMD-MGCM to study
the NO nightglow detected by SPICAM [7]. This is
a process specially suitable to be studied with a
GCM like ours, because it implies a strong coupling
between radiation, dynamics and chemistry. For this
purpose, Nitrogen chemistry has to be included in
the model.
An ongoing effort is devoted to the
implementation of a first ionospheric module into
the GCM [11], which will allow us to extend our
results to further measurements from Mars Express,
like the variability of the altitude and peak of the
electron densities.References: [1] Forget, F. et al. (1999), JGR 104,
24155-24175. [2] López-Valverde, M.A. and M.
López-Puertas (2001), ESA technical report. [3]
González-Galindo, F. et al. (2005), JGR 110,
DOI:10.1029/2004JE002312. [4] Angelats i Coll, M. et
al., (2005) GRL 32, DOI:10.1029/2004GL021368. [5]
Keating, G. et al (1998), Science 209 1672-1676. [6]
Bougher, S. et al. (2006), GRL 33,
DOI:10.1029/2005GL024059. [7] Bertaux, J.L. et al.
(2005), Science 307 , 566-569. [8] Bougher et al. (2000),
JGR 105, 17669-17692. [9] González-Galindo et al., in
preparation. [10] Forget et al., this issue., [11] Gilli, G. et
al., this issue
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
MAPPING OF PLAINS VOLCANISM IN TEMPE TERRA, MARS: NEW OBSERVATIONS WITH
POST-VIKING DATA. E. Hauber1, P. Bro
1, J. Bleacher
2, D. Williams
3, R. Greeley
3.
1Institute of Planetary
Research, DLR, Rutherfordstr. 2, 12489 Berlin, Germany.2NASA/GSFC, Greenbelt, MD, 20771, USA.
3ASU,
Tempe, AZ, 85287-1404; USA. [email protected]
Introduction: Viking Orbiter (VO) imagesrevealed that the western Tempe Terra region (TT)
on Mars (located in the NE portion of Tharsis)
displays various surface features that are indicative
of basaltic volcanism. The morphologic evidence
includes coalescing low shields, fissure vents, pit
craters, steep cones, and lava flows [1]. This
assemblage of volcanic landforms is very similar to
that of the Snake River Plains (Idaho; USA), where
Greeley [2] defined the term plains volcanism for a
style of volcanism that is intermediate between
Hawaiian shields and flood basalts. The most
detailed study of TT since Plescia`s work [1] is a
USGS geologic map [3], still entirely based on VOimages. Recent studies began to incorporate
accurate topographic data (MOLA PEDR`s) and
high-resolution MOC images [e.g., 4-6]. A signifi-
cant improvement in the available data base is now
provided by HRSC data [7], which fully cover the
region, THEMIS data, and the increasing number of
extremely high-resolution HiRISE images. We show
detailed topographic investigations of low shields
and present our HRSC-based mapping of volcanic
landforms in Tempe Terra.
Summary of Observations: As suggested earlier
[1], TT is widely covered by volcanic material. The
extremely low flank slopes of the shield volcanoes(Fig. 1) suggest a very low viscosity of lavas. This
could be a result of high eruption temperatures, high
effusion rates, or a low Si- and a high Mg-content
along with a possibly high Fe-content of the lavas
[6]. HRSC-based mapping (Fig. 2) shows an even
denser pattern of vents, which are controlled by a
pre-existing, NE-trending tectonic pattern, and
coalescing shields than it was obvious in Viking-
based mapping [3]. However, many kipukas are
present and low shields and associated lava plains
become smaller and more isolated towards the SE.
Both facts point towards a relatively small total
thickness of the volcanic cover in TT, which might
represent a thin late-stage veneer of basalt above
much older and tectonically deformed basement.
Our observations confirm previous reports of
volcanism in TT. In addition, we describe new
features (cinder cones, sinuous rilles) that have not
been discussed in the context of TT before. In
conclusion, we confirm the notion that the Snake
River Plains [2] are the best terrestrial analog to
Martian low shields and their associated landforms.References: [1] Plescia, J. (1981), Icarus 45, 586-601.
[2] Greeley, R. (1982) JGR 87 , 2705-2712. [3] Moore,
H. J. (2001) U.S.G.S. Geol. Inv. Series I-2727.
[4] Sakimoto, S. et al. (2003b) Sixth Int. Mars Conf., Abs.#3197. [5] Wong, M. et al. (2001) LPS XXXII , Abs.
#1563. [6] Hauber, E. (2007) 7 th Int. Conf. on Mars, Abs.
#3287. [7] Jaumann, R. et al. (2007) PSS 55, 928-952.
Figure 1. Low shield in Tempe Terra (~36°N/272°E).
Note the almost perfectly symmetrical shape, which is
enhanced by the 25 m contour lines. The extremely low
relief (~250 m height, diameter >20 km) generates little
shading in the images. This is similar to Icelandic low
shields (right), which are only recognizable in Landsat
images due to the snow cover at higher elevation.
Figure 2. Detail of VO-based map of volcanic features in
Tempe Terra. Note that many coalescing low shields are
clearly visible from the 20 m contour line arrangement
(derived from gridded MOLA data). Vents are trendingNE and are controlled by the regional tectonic trend.
Impact craters and their ejecta are shown in green.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Mapping of Spectral Units using HRSC Color Data. T. B. McCord1. J.-Ph. Combe
1,
1Bear Fight Center, 22
Fiddler’s Road, Winthrop, WA, 98862, USA. [email protected]
The general focus
HRSC has produced a very large data set with
four spectral bands specifically engineered toprovide color data across the extended visible
portion of the spectrum. These data have been
characterized and initial example attempts described
to utilize them for spectral and compositional
analysis of the Mars surface [1]. A general result
was that the Mars surface in the regions studied
could be modeled mostly using only a few spectral
components, namely red rock = iron oxide-rich
material, dark rock = unoxidized basalt, fresh ice (at
the poles), and a shade/shadow component, or their
mixtures [1]. Yet, OMEGA has produced spectra
indicating the existence of additional compositional
(e.g., salts, phyllosilicates, hematite) and spectralcomponents using much higher spectral resolution
IR spectra. A major advantage of the HRSC color
data over the OMEGA observations is their much
higher spatial resolution and wide area coverage
plus topographic context. An attractive possibility is
to use the OMEGA spectra to identify
compositional components and the HRSC color
images to map these components at a much higher
spatial scale over a wider region. Thus, we have
initiated several studies to attempt this approach.
HRSC spectral data analysis
The previous study [1] showed that the HRSC 4-
color data could be treated as a 4-component vectorin a 4-space for each pixel. One can visualize the
data cloud for each scene in several ways and treat
the data sets using a technique called Spectral
Mixture Analysis (SMA) [2] to calculate the
proportion of each identified spectral component
within each pixel. The previous analysis was
unsuccessful at finding convincing evidence of
unaccounted-for endmember materials beyond those
stated above in the few scenes analyzed. This may
be partly the result of effects of HRSC data
compression, registration errors among the color
channel data, and other undesirable effects. Yet,
there were signs of spectral contributions at the level
of the “noise” in the data that could be indications of
unaccounted-for effects of surface roughness/shade
and other spectral components associated with
unidentified materials, for example as shown in Fig.
21 from [1]. We continued the search for
expressions of other compositional components in
the HRSC color data and report here on our results
so far.
Echus Plateau
Because of its relevance to a larger study, we
attempt to model this region using the SMA
technique [3, 4]. Our first attempt showed that onemust account more precisely for the shade/shadow
component and remove it before developing the
final model, as was reported also in [1]. This effect
was studied separately and a successful model
developed [Combe et al., this conference]. Afterapplying this model and removing the shade/shadow
component, we find that this region can be modeled
very well using only the red rock and dark rock
components and their mixtures, consistent with the
earlier study; described by Combe et al., this
conference.
Search for salts
Another study underway is to treat regions that
are reported from OMEGA data to contain outcrops
of salts (magnesium sulfate) and other materials.
Some salt deposits are clearly visible in the HRSC
panchromatic data as bright, well-defined units.However, when we modeled the first of these areas
(the Juventae Chasma unit described by [5] and
using Orbit 243 HRSC data), we found that there is
no unique spectral characteristic of this deposit in
the HRSC color data that allows distinguishing them
from mixtures of red rock with dark rock [described
by Wendt et al., this conference]. This may be
because the salt spectra are featureless (flat) in the
visible spectral region and are multiple scattering
and therefore easily take up the color of any colored
contaminate—dark rock in this first example. We
are treating more regions and will report on our
further experience.Results so far
Our attempt to detect spectral evidence of Mars
surface compositional components beyond what was
reported earlier [1] have so far been unsuccessful,
but we are continuing our effort. In the process, we
have further characterized the HRSC color data and
are producing spectral/compositional maps of
interesting Mars regions for at least the major
spectral/compositional components. Further, we
have increased our ability to detect, determine and
map surface roughness using the shade component
from the SMA technique. This work is continuing
and further recent results will be reported with
details of the approach used.
References: [1] McCord, T. B. et al. (2007), JGR 112,
doi:10.1029/2006JE002769. [2] Adams and Gillespie
(2006), Cambridge U. Press. [3] Combe, J.-Ph. et al:
Analysis of OMEGA / Mars Express data hyperspectral
data using a Multiple-Endmember Linear Spectral
Unmixing Model (MELSUM): Methodology and first
results. Submitted to Planetary and Space Sciences. [4]
Combe et al., (2006) LPSC XXXVII , Abs. #2010. [5]
Gendrin, A. et al. (2006), LPSC XXXVII , Abs. #1872.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
MARS EXPRESS SPICAM UV DATA: RESULTS ON THE UPPER ATMOSPHERE OF MARS. C.
Simon1, O. Witasse
1, F. Leblanc
2, J. Lilensten
3, and J.-L. Bertaux
2.
1ESA/ESTEC/RSSD Noordwijk, The
Netherlands2
CNRS/aeronomie Verrieres Le Buisson, France3LPG Grenoble, France
This paper is intended to describe some results
obtained by the SPICAM experiment aboard MarsExpress. Nadir and limb data are analysed in order
to obtain information on the aeronomy of Mars. An
updated Point Spread Function (PSF) is used to
estimate the intensities of the airglow. This allows a
better calculation of the vertical profiles
of the atmospheric emissions. Results are compared
with numerical simulations.We focus on the following areas:
- Dayside carbon monoxide 'Cameron Bands' and
CO2+
emissions: seasonal behavior,
- search for the O+
emission at 247 nm,
- atomic oxygen emissions at 297 and 135 nm, and
- nightside auroral emissions.
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AN EMERGENT, NEW PARADIGM FOR MARS GEOLOGY. M. C. Malin and K. S. Edgett, Malin SpaceScience Systems, P.O. Box 910148, San Diego, CA 92191-0148 USA ([email protected]).
The planet Mars revealed by the Mars Global Sur-veyor (MGS) Mars Orbiter Camera (MOC) is not theMars we all thought we knew from the Viking and
Mariner missions. One of the most fundamental (andusually unstated) assumptions about martian geologyis the notion that the planet’s heavily cratered terrain islike that of the Moon—a megabreccia of primordialcrust, perhaps consisting of the same lithologythroughout (e.g., lunar anorthosite) . The additionalunstated assumption is that “heavily cratered terrain”formed on the surface of Mars, and then “stuff hap-pened” to this terrain—such as mantling, volcanicplains formation, and processes involving running wa-ter—to provide the configuration of landforms that wesee today. Nearly all pre-MGS efforts to classify, quan-tify, map, and date these surfaces on the basis of appar-ent stratigraphic relationships and crater counts have
made this basic, though simplifying, assumption.MOC images reveal that the upper crust, every-
where that it is exposed in outcrop form, is layered(e.g., Figs. 1, 2). In retrospect, the Mariner 9 andViking Orbiter images also demonstrate this fact, oftenin dramatic ways that did not make sense until now.Additionally, the surface properties of Mars, such asthe regional albedo patterns and the distribution of dunes, yardangs, rocks, and dust, are all manifestationsof this layered upper crust and subsequent weatheringand redistribution of its materials.
MOC images show that there are different types of layers, some with different albedo and some with differ-ent resistance to erosion. They also hint that some of the layers must be of sedimentary origin, and mighthave a regional extent that implies that processes oc-curred on ancient Mars that are completely unlike theprocesses that occur there today. Layered rocks (espe-cially sedimentary rocks) on Earth tell tales of entireseas and mountain ranges that have formed, evolved,and vanished.
The presence of a layered crust suggests that planetMars was once unlike anything that anyone has everdescribed. This Mars existed at a time when impactcrater formation was still frequent, and probably existedat a time that predates all of the major volcanic andtectonic features of the Tharsis/Syria rise. This is anearly Mars that pre-dates all of the landforms previ-ously attributed to “early Mars”—for example, valleynetworks—and the surface that has been described priorto MGS is simply that found within the final few chap-ters in a diverse and previously unrecognized martianhistory. This Mars is just barely accessible to space-craft, and will likely require careful exploration by hu-man geologists operating in the field for many decadesto fully reveal and appreciate the complex and rich his-tory of this terrestrial planet.
Figure 1. Layered outcrop exposed in wall of eroded im-pact crater that is superposed on a fretted terrain valley innorthern Arabia Terra. Image located near 38.3°N,320.8°W. North is approximately up, illumination is fromthe right. Subframe of MOC SPO2 image 46502.
Figure 2. Layers expressed as terraces with three differentcraters in three different states of exposure or exhumationin the heavily cratered terrain of central Arabia Terra. Im-age located near 19.2°N, 353.6°W. North is approxi-mately up, illumination is from the left. Subframe of MOCSPO2 image 53403.
Fifth International Conference on Mars 6027.pdf
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2. The Planet Mars 1
Making Tracks on Mars New Mexico Museum of Natural History & Science LodeStar Astronomy Center•
2. The Planet Mars
A . Mars and New Mexico
There are connections between Mars and New Mexico
1. New Mexico has a long history of participation in the development of technologyused in Mars missions, from Goddard’s rocket experiments at White Sands tomodern work on robotics, instruments, and air bags at Los Alamos NationalLaboratory (LANL) and Sandia National Laboratories (SNL).
2. Scientists and engineers from many different institutions in New Mexico havebeen funded by NASA to do fundamental Mars research or to work as members
of Mars mission teams. These scientists and engineersinclude Dr. Larry S. Crumpler, research curator, New
Mexico Museum of Natural History & Science(NMMNHS), science team member for the Mars Explo-
ration Rover mission and instrument team member forthe Mars Odyssey mission; Dr. Robert Reedy, geochem-ist from Los Alamos National Laboratory (LANL) (cur-rently at the University of New Mexico (UNM)), instru-ment team member for the Mars Odyssey mission; Dr.Penny Boston, biologist from New Mexico Tech, work-ing with NASA astrobiology to research extremophilesin New Mexico caves; and Dr. Carl Agee, Dr. HortonNewsom, and Dr. Barbara Cohen, a newly named
member of the MER science team as of 2006, all from the Institute of Meteoriticsat UNM, funded to study meteorites from Mars. This list includes only a fewNew Mexican research scientists currently funded by NASA to study Mars orMars-related science. There are many others statewide.
3. Geologically, New Mexico and Mars have many similarities, and our state canbe used to understand some of the physical
properties and processes operating on Mars. Inmany ways, the surface of Mars is a lot like NewMexico: it is geologically diverse, with plains,ridges, buttes, mesas, volcanoes, canyons, land-
slides, and arroyo-like channels. It displays anarid environment with evidence of past water.
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2. The Planet Mars 2
Making Tracks on Mars New Mexico Museum of Natural History & Science LodeStar Astronomy Center•
FAST FACTS: A Mars day is 24 hours,
37 minutes and is called a sol. A Mars
year is 687 Earth days. Gravity on
Mars is a about one-third that of
Earth.
B. Mars : Physical Propert ies
How far away from Earth is it?
Mars is the fourth planet from the sun, located between Earth and Jupiter. Because
the distance varies depending upon their locations in their relative orbits, Marsranges from 35,000,000 miles (56,000,000 km) to 249,000,000 miles (399,000,000km) from Earth. If you could travel from Earth at a constant rate of 60 miles perhour (about 97 km/hr), it would take 66 and a half years to get to Mars when it is atits closest approach to our planet. In August of 2003, the two planets were theclosest they have been in about 60,000 years, which is the reason both NASA and
the European Space Agency (ESA) launched missions to Mars during that time.
How big is Mars?
Mars is 4,200 miles (6,800 km) in diameter at the equator. That means Mars isabout one-half the size of Earth and twice the size of Earth’s moon. Although theplanet is smaller than Earth, Mars has approximately the same land area as Earthbecause there are no oceans. Gravity is related to the mass of a planet. Gravity onMars is about one-third that of Earth (which means on Mars you would weigh about
a third your weight on Earth and you could dunk a basketball in a basket threetimes higher than on Earth). Our current knowledge suggests that Mars does nothave a magnetic field similar to Earth’s. Other physical properties of Mars, and acomparison with those of Earth, are listed in Table 1.
Does Mars have a moon?
Mars has two natural satellites, or moons, but they don’t look like Earth’s moon.
Mars moons are named Phobos (Fear) and Deimos (Panic) (from Greek mythology,the sons of “Ares” or “Mars,” the god of war. They are very small and irregularlyshaped. In fact, they are frequently described as looking like potatoes. Phobos isonly about 12 miles (20 km) in diameter and Deimos is about 7 miles (12 km) indiameter. They have dusty surfaces with impact craters of many sizes, and they arebelieved to be captured asteroids in orbit around Mars.
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2. The Planet Mars 3
Making Tracks on Mars New Mexico Museum of Natural History & Science LodeStar Astronomy Center•
Table 1.
Mars: Basic Information
Mars Earth
Diameter at equator 6,794 km 12,756 km
Dry surface area 144X106 km2 148X106 km2
Gravity 0.38 G (3.71 m s-2) 1.0 G (9.78 m s-2)
Mass 6.4185X1023 kg 5.9736X1024 kg
Volume 1.632X1011 km3 (0.15 of Earth) 1.083X1012 km3
Density (water=1) 3.93 g cm-3 5.52 g cm-3
Day 24 hr, 37 min 24 hr
Year 687 Earth days 365 days
Atmosphere 95% Carbon Dioxide 76% Nitrogen
2.7% Nitrogen 21% Oxygen
1.6% Argon 0.9% Argon
0.13% Oxygen 0.03% CO2
Mean Surface Pressure 6.36 millibars 1.0 bars
Mean Temperature
(at equator) -70° F 80° F
Albedo 0.16 0.29
Axial Inclination 25.19° 23.44°
Natural satellites
(moons) Phobos Moon
(Diameter: 20x23x28 km) (Diameter:
3,476 km)
Deimos
(Diameter: 10x12x16 km)
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2. The Planet Mars 4
Making Tracks on Mars New Mexico Museum of Natural History & Science LodeStar Astronomy Center•
FAST FACT: The Mars volcano
named Olympus Mons is the
largest volcano in the solar system.
Its base would cover the entire
state of New Mexico.
C. The Surface of Mars
Why is Mars called the Red Planet?
Mars is called the Red Planet because it appears to glow a ruddy red color in the
night sky as viewed from Earth. The surface of Mars is red, actually reddish-brown;the color is caused by rusting (or oxidation) of iron minerals in the rocks and dirt atthe surface. Planet-wide dust storms move the red dust around to coat the entiresurface and make the sky of Mars pinkish in color.
Is Mars similar to Earth?
Mars is a geologically diverse planet like our own planet Earth. The surface hasbeen affected by volcanism, faulting, impact cratering, and the action of wind, wa-ter, and ice.
Topographically, there is a major difference between Earth and Mars. The surfaceof Mars can be divided into two areas: high-standing ancient cratered highlands andlow-lying plains. The cratered highlands include most of the southern hemisphere of Mars, while the low-lying plains form most of the northern hemisphere of Mars. The
ancient cratered terrain is similar to the cratered highland areas of the Earth’smoon. A key difference is that Martian cratered highlands are more eroded. Impactcraters smaller than 10 km (6 miles) across are nearly all eroded away; only largercraters remain in the Martian cratered highlands. We do not know what caused thiswidespread erosion; and, although many theories have been proposed, we do notknow what caused the distinctly different southern and northern hemispheres of Mars.
Are there volcanoes on Mars?
There are about 20 very large Martian volcanoes (greater than 300 miles or 500 kmin diameter) and numerous smaller ones. Volcanoes occur mainly in three regionson the planet. One area is a large topographic bulge known as the Tharsis region. A
smaller bulge on the other side of the planet, known as the Elysium region, andareas in the southern cratered highlands also include big volcanoes.
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2. The Planet Mars 5
Making Tracks on Mars New Mexico Museum of Natural History & Science LodeStar Astronomy Center•
Olympus Mons, in the Tharsis region, is the largest volcano in the solar system. Itis over 370 miles (600 km) across, and its base would cover the entire state of NewMexico. Its summit crater, or caldera, is as wide as the distance from Albuquerqueto Santa Fe. In many spacecraft images, the upper part of Olympus Mons appearssharper and clearer, and the lower part looks foggy. This is because the top is stick-
ing up above most of the dusty atmosphere. It is about 13 miles or 70,000 feet high(22 km), about twice the height of Mt. Everest. Olympus Mons was named after thehome of the Greek gods of mythology, Mount Olympus. Astronomers once called itNix Olympica (the snows of Olympus) because they could see a white spot at itslocation through Earth-based telescopes; however, this white spot was probablyclouds formed at the summit, not snow or frost. Amazingly, although Martian volca-noes are very large, they appear to be typical “shield volcanoes,” built up by incre-mental lava flows, similar to those that form the Hawaiian Islands on Earth or San
Felipe shield volcano in New Mexico. Scientists believe that all of Mars’ volcanoesare very old and no longer active.
Are there marsquakes?
Volcanic eruptions and marsquakes have formed some of the deepest valleys, big-gest plains, and tallest volcanoes in the solar system on Mars.
Surface ridges, cracks, and fractures provide evidence for the existence, at least in
the past, of marsquakes. The most spectacular example in the solar system is thelarge system of canyons, known as Vallis Marineris, located near the equator andon the eastern flanks of the Tharsis bulge. The origin of Vallis Marineris is not verywell understood, but faulting was important. There is no evidence for Earth-likeplate tectonics on Mars, so this canyon is not a plate boundary. It is actually a
geologic rift, like the Rio Grande rift, where the crust of the planet has thinned andpulled apart. Our Grand Canyonwould be a small tributary canyonto Vallis Marineris, which extends
for more than 2,500 miles (4,000km), and would reach from New
York City to Los Angeles if it wereon Earth. At its greatest width,
Vallis Marineris is 365 miles (600km) across and up to 4 miles (7 km)deep. Faulting may have createdmost of the initial canyon, but parts
of the walls have collapsed in hugelandslides that continue to widenthe canyon. Vallis Marineris wasnamed after the Mariner 9 space-craft, the first orbiter spacecraft toimage the entire surface of Mars.
Comparison of Olympus Mons and New Mexico
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2. The Planet Mars 6
Making Tracks on Mars New Mexico Museum of Natural History & Science LodeStar Astronomy Center•
Self-portrait of the rover named Spirit, Sol 111
Is there water on Mars?
At the present time, water is permanently found on Mars only at the poles. Here itis so cold that water ice can last at the surface. Long-lived ice sheets or “caps” atboth poles consist of carbon dioxide ice (also called dry ice), water ice (a small
amount), and layered sediments. The Mars polar caps decrease and increase in area
seasonally, just like the polar ice caps of Earth. There is also evidence that addi-tional water now exists as frozen permafrost or subsurface ice in the middle to highlatitudes.
Mars’ atmospheric pressure today is too low to support the existence of liquid waterat the surface. However, there is evidence of abundant past liquid water at thesurface, because we see old water-cut channels and obvious flood deposits. In fact,
valleys and channels similar to water-eroded valleys on Earth occur all over Mars.The floors of some valleys look like river valleys with many channels. Other valleysare similar to arroyos in New Mexico formed by runoff from slopes. Some large
valleys converge and appear to drain into low basins. Sediment appears to have
been deposited in some of these low basins and in some craters. The valleys that aremost like river channels occur in the eroded ancient cratered highlands, which
suggests that Mars’ climate was wet very early in the history of the planet. So muchwater has apparently flowed on Mars that some geologists have argued that largebodies of water could have covered vast areas of the lowland plains. Finding outwhat happened to the once-abundant water and whether there were ever oceans,lakes, or rivers on Mars was one of the goals of the Mars Exploration Rover mission.
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2. The Planet Mars 7
Making Tracks on Mars New Mexico Museum of Natural History & Science LodeStar Astronomy Center•
D. The Atmosphere of Mars
Could we breath Mars air?
Mars’ atmosphere is very thin and unbreathable, like the atmosphere of Earth
about 200,000 feet (38 miles) straight up. The surface atmospheric pressure onMars is only 6 mbar versus 1,000 mbar (1 bar) on Earth. The atmosphere is about95 percent carbon dioxide with a little nitrogen, argon, and water. It is also dusty,which makes the sky of Mars pinkish-yellow in color. We worry about possibleglobal climate change on Earth; the climate of Mars appears to have changed radi-cally over its geologic history. Scientists believe that Mars once had a thicker,
warmer atmosphere but don’t yet know why and how Mars changed.
What is the temperature on Mars?
Currently Mars is cold! The average temperature on Mars is minus 81 degreesFahrenheit (minus 63 degrees Centigrade). Generally the temperature ranges froma high of freezing (32 degrees F or 0 C) on a summer day at the equator down tominus 250 degrees F (minus 157 degrees C). Because the air is so thin, if you couldstand at the equator, the surface temperature would change from 70 degrees F atyour feet to 32 degrees F at the top of your head! And Martian air is dry, even com-pared to New Mexico. It contains only about 1/1,000 as much water as our air. Eventhis small amount can condense out, forming clouds high in the Mars atmosphere oraround the slopes of the big Martian volcanoes. Local patches of early-morning fog
can form in the valleys of Mars.
Is there weather on Mars?
Mars has weather, of a kind, and seasons. The rotation axis of Mars is tilted fromthe vertical axis at about the same degree as Earth; therefore, it has four seasons(but remember, a Mars year is twice as long as Earth’s; therefore, each Mars seasonis twice as long).
Currently, Mars weather consists of wind, dust storms, dust devils, occasionalmorning frost (made of dry ice or frozen carbon dioxide), and clouds, especiallyaround the highest peaks. Cold fog often forms around the polar caps. Although theclimate has gotten drier, Mars
still has an Earth-like systemof weather, with low-pressuresystems similar to the onesthat cause storms on Earth,
but no snow or rain sinceliquid water cannot exist atthe surface of Mars under itscurrent atmospheric pressure.
Martian clouds as seen by the Opportunity rover from Endurance Crater on Sol 291
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2. The Planet Mars 8
Making Tracks on Mars New Mexico Museum of Natural History & Science LodeStar Astronomy Center•
Mars is windy, but even though the wind speeds range from 3 to 70 mph (5 to 113kph), the low atmospheric pressure means that only small particles (dust ratherthan sand) can be moved by the wind. Daily winds circulate air from high to lowelevation (off volcanoes or mesas) just as in New Mexico with its afternoon winds.Dust devils, like those that form on a summer day in New Mexico, are very common
on Mars. Dark dust devil tracks can be seen in many areas of the surface, and largedune fields are common. Global winds on Mars produce large seasonal dust storms.
A third of the carbon dioxide in the atmosphere of Mars freezes each winter at thepolar ice caps. This causes tremendous swings in atmospheric pressure. Thesepressure swings are partly responsible for the dust storms that can nearly obscurethe surface of the planet in the spring.
FAST FACT: The Viking Mission landers
sent back data on the weather of Mars,
which includes wind, dust storms, dust
devils, morning frost, clouds, and fog. The
atmosphere is thin and the air is very
cold.
A dust devil (top) viewed by Spirit from the summit
of Columbia Hills.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
MARS ORIGINS MISSION F. Westall1, G. Klingelhöfer
2, and the Mars Origins Scientific Team
3.
1Centre de Biophysique
Moléculaire, CNRS, Rue Charles Sadron, 45071 Orléans cedex 2, France.2Johannes Gutenberg-Universität Mainz, Institut für
Anorganische und Analytische Chemie, Staudinger Weg 9, D-55099 Mainz, Germany. westall@cnrs-orléans.fr ;
The Mars Origins Mission [1] is an in situ mission to theNoachian terrains of Mars with the following scientific
objectives :
1. To characterise the very early geological evolution of
Mars and the context in which life potentially arose;
2. To search for traces of the transition from a prebiotic
world to life;
3. To trace the early evolution of life and its fate as
conditions on Mars changed.
This mission will provide information on the first billion
years of inner planet evolution and the appearance of life,
i.e. critical information that is lacking on Earth.
Similarities in the mode of formation of the terrestrialplanets, Earth, Mars and Venus, and in their sizes andorbital positions suggest similar environmental conditionson these planets early in their histories. All three planetshad a CO2 rich atmosphere enabling the presence of liquidwater on their surfaces, as well as an abundance of prebioticorganic molecules. With the addition of other ingredientsessential for life (H, O, N, P, and S) and a source of energy(geothermal, geochemical or solar), Earth, Mars and Venuscould have supported the independent appearance of life.Environmental conditions on all three planets havechanged throughout geological time. The Earth developedplate tectonic processes and maintained its atmosphere,
which was entirely altered by the products of abundant life(i.e. oxygen). Mars lost its magnetic dynamo, lost a largepart of its atmosphere, and became dry and frozen planet,while Venus suffered a “runaway greenhouse”. On Earth,vigorous plate tectonic activity has largely erased the firstbillion years of geological and palaeontological history.Although rocks as old as ~4 billion years exist on theEarth, they are so heavily altered and deformed that it isdifficult to obtain information about the earlyenvironmental conditions. By 3.5 Ga, the time that oldestwell-preserved rocks formed, the geological, environmentaland palaeontological records that they contain pertain to arelatively geologically evolved planet inhabited byrelatively evolved life forms.
Mars, the smallest of the three planets cooled morequickly and froze tectonically, probably between 4.2-3.8Ga. This is of fundamental importance to our scienceobjectives because, in its early Noachian terrains, Marsstill retains the record of the first billion years of evolutionof the terrestrial planets that has been erased on Earth. Thegeological context of early Mars therefore forms thebackdrop for investigating the “missing link”, the earlygeological evolution of and the origin of life on theterrestrial planets.
Moreover, irrespective of whether life appeared on Mars,and irrespective of whether traces of life occur within reachof the rover at the landing site, Mars remains a
calibrationary dipstick in terms of testing null hypothesesrelated to traces of life. Lacking vigorous tectonic recyclingover the last >3.5 Ga, crust dating back to the time of thedifferentiation of the planet still exists at the surface, orclose to the surface.
The scientific objectives will be addressed by a ~40 kg
payload of rover-based instrumentation inherited partly
from already existing (but improved) technology of
Beagle 2, ExoMars, and MER, and partly on completely
new instrumentation. A specific goal will be rock dating
by rock/mineral isotopic analysis (with an accuracy of 100
My). Dating is fundamental to the objectives (e.g. dating
the cessation of the martian dynamo and being able toselect and cache of the most relevant Noachian-aged
samples in preparation for a Mars Sample return mission).
Other new instruments include a magnetometer and an
electron gun for cathodoluminescence. The instrument
suite includes: a panoramic camera; Marsfly (flying
geologist); a close-up imager/microscope; a
magnetometer; Raman LIBS; K-Ar and 40Ar-39Ar
dating; a multispectral microscope; GC-MS;
Mössbauer/XRF; and XRD/XRF. Sample acquisition will
use a 60 cm drill.
In view of preparations for a future Mars Sample Return
mission (2020), our Mars Origins Mission will be
extremely timely for testing systems that will be necessaryfor such a mission, such as a caching system for samples
of the highest scientific interest (origins), a
communications orbiter; lander insertion from orbit; soft
landing; precision landing (10 km ellipse); rover (rough
terrain capabilities); alternative energy sources (RTGs for
longevity, night operations); direct command of the rover
with provision of communications from orbit; and data
retrieval from orbit. The spacecraft can be launched using
a Soyuz Fregat launcher in a direct hyperbolic transfer to
Mars.
References:[1] Westall, F and Klingelhöfer, G., 2007.
Mars Origins Mission. ESA.Cosmic Vision proposal. [2]Westall, F., 2005. Early Life on Earth and Analogies to
Mars,in T. Tokano (Ed.) Water on Mars and Life, pp. 45–
64
3.D. Pullan, T. Zegers, R. Arvidson, B. Hofmann, A. Coradin, J.-L. Josset,
A. Griffiths, R. Jaumann, I. Wright, F. Rull, E. Jessberger, D. Tallboys, L.Marinageli, M. Trieloff, P. Ehrenfreunde, L. Becker, F. Gössmann, C.
Ramboz, C. Briois, C. D’Uston, B. Weiss, C. Schröder, N. Arndt, M.Grässlin, R. Laufer,H.-P. Röser, O. Zeile, S. Gorevan, J. Bada, M.
Madsen, F. Raulin, A. Brack, J.-P. Bibring, M. Sims, R. Morris, A.Hofmann, C. Cockell, D. Breuer, H. Edwards, J. Parnell, C. Sotin, F.
Costard, N. Mangold, M. Toplis, D. Lentink, C. Passchier.
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THE MARS OXIDANT INSTRUMENT R. C. Quinn1, A. P. Zent
2, F. J. Grunthaner
3, P. Ehrenfreund
4 1SETI
Institute, NASA Ames Research Center, MS 239-4, Moffett Field, CA 94035, USA,2 NASA Ames Research
Center, Moffett Field, CA, USA,3 NASA Jet Propulsion Laboratory, Pasadena, CA, USA,
4Astrobiology Group,
Leiden Institute of Chemistry, Leiden University, NL. [email protected]
Urey: Mars Organic and Oxidant De-tector, has been selected for the Pasteur
payload in the European Space Agency’s
(ESA’s) ExoMars rover mission. Part of
the Urey instrument suite, the Mars Oxi-
dant Instrument (MOI) uses a chemical
sensor array to characterize the chemical
processes that modify organic compounds
in the Martian environment. MOI pro-
vides a simple yet robust method of as-
sessing the oxidant characteristics of
sample material and evaluating how oxi-
dation reactions may have altered the
original organic components. MOI meas-ures the reaction rates of films that have
different sensitivities to particular types
of oxidants expected to be present in the
Mars surface environment. By controlling
the temperature of these films and their
exposure to dust, ultraviolet light and wa-
ter vapor, MOI will evaluate organic deg-
radation pathways that may take place at
sampled localities on Mars. These data
will provide important insights into the
observed organic matter inventory and the
potential for survival of various classes of
organic compounds under Martian envi-ronmental conditions. MOI components
(Figure 1) have been designed for modu-
larity, ease of assembly, and the ability to
unambiguously measure the extent and
character of chemical reaction between
the sensing films and atmospheric oxi-
dants, dust, soil and UV. The MOI films
are chosen to emulate biotic and abiotic
materials that may be, or may have been
present, in the field site environment.
Analysis of the combined Urey results
(MOD /μCE/MOI) will allow mecha-
nisms that are modifying organic compo-nents in the Martian environment to be
characterized. In combination with a di-
rect search for organic compounds by
other payload instruments, MOI will al-
low characterization of the carbon cycle.
These results will be especially critical in
the case of a failure to detect organic
components by other instruments, or the
detection of modified organic intermedi-
ates. We propose that a key to under-
standing carbon chemistry on Mars lies
not only in identifying soil oxidants, but
also in characterizing the dominant reac-tion mechanisms and kinetics of oxidative
processes that are occurring on the planet.
These processes may have decomposed or
substantially modified any organic mate-
rial that might have survived from an
early biotic period.
Figure 1. Six MOI chemical sensors.
Acknowledgements: The authors thank
Jeff Bada (University of California at San
Diego) and the entire Urey team.
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European Space AgencyEuropean Mars Science and Exploration Conference: Mars Express & ExoMarsESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THE MARS SCIENCE LABORATORY (MSL) RADIATION ASSESSMENT DETECTOR (RAD) ANDIMPLICATIONS FOR (IRAS) ON EXOMARS
R. F. Wimmer-Schweingruber1, D. M. Hassler2, O. Kortmann1, E. Böhm1, C. Martin1, R. Beaujean1, S.
Burmeister1, A. Posner3,4, G. Reitz5, and the MSL/RAD Team
1
Institute for Experimental and Applied Physics, University of Kiel, Leibnizstr. 11, 24098 Kiel, Germany.2Southwest Research Institute, 1050 Walnut St, Ste 400, Boulder CO 80302, USA. 3Southwest Research
Institute, 6220 Culebra Rd, San Antonio, Texas 78238-5166, USA, 4Currently at Science Mission Directorate,
NASA HQ, 300 E St. SW, Washington, DC 20546, USA. 5German Aerospace Center (DLR), AerospaceMedicine, Linder Höhe, 51147 Cologne, [email protected]
The Radiation Assessment Detector (RAD) onNASA's Mars Science Laboratory mission is beingbuilt to characterize the broad-spectrum of thesurface radiation environment, including galacticcosmic radiation, solar proton events, and secondaryneutrons.
This overarching mission goal is met by RADsscience objectives 1-5: 1.)Characterize the energeticparticle spectrum incident at the surface of Mars,including direct and indirect radiation created in theatmosphere and regolith. 2.)Validate Marsatmospheric transmission models and radiationtransport codes.3.)Determine the radiation Dose rate and EquivalentDose rate for humans on the Martian surface.
A pathfinder model with flight-like properties has
been tested at BNL, iThemba, CERN/CERF, and
using various radioactive sources to demonstrate the
measurement capabilities required by its science
objectives. We will present first calibration results
and compare them with GEANT4 simulations.
The neutron-gamma discrimination can be achievedin a statistical manner using a combination of
different scintillators1 and will also presented. Finally, we will discuss implications for the IonizingRAdiation Sensor (IRAS) for ESA's ExoMarsmission.
References:1
E. Böhm, A. Kharytonov, andR.F.Wimmer-Schweingruber, A&A preprintdoi: http://dx.doi.org/10.1051/00046361:2007726
Figure 1: MSL/RAD sicene goals.
Figure 2: Comparison between measurement (black) andsimulation (red and green) for iron fragments.