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.






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.






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.

[email protected]

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.






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.






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






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],

[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.



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



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].

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(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.,

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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).






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






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.

[email protected]

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.

mailto:[email protected]









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






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.






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

[email protected]

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.






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.

[email protected]

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






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.






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.






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






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.

[email protected].

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.






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.






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.






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.

[email protected]

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






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).






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.

[email protected]

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.






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.






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.






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.






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)






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.






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.

[email protected]

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






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.






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.






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



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



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.

unar and Planetary Science XXXI 1058.pdf






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.






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.






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.



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

Lunar and Planetary Science XXXIV (2003) 1134.pdf



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






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).






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






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.






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

.






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






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



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)






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








MATHEMATICAL ASPECTS ASSOCIATED TO THE GENERATION OF IMPACT CRATERS ON

MARS J. C. Echaurren1.

1Codelco Chile Chuquicamata, North Division, Pasaje Lince 976, Calama, Chile.

[email protected]

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.






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

[email protected]

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






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.






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.






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






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.






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.






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.






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.






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..

[email protected]

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.






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.






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

[email protected]

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.






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.






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

[email protected]

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






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.






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

4500

G a n g i s

G a n g i s

G a n g i s

G a n g i s

G a n g i s

E o s / C a

p r i

A r s i n o

e s

A u r e u m

A u r e u m A r

a m I a n i

I a n i

T h i c k n e s s [ m ]






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.






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






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.






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.






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.



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



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.





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.





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)





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.





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





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.





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





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.






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 ;

[email protected]

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.






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.



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.

esa mars research abstracts part 1

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