Experimental study of the sublimation of ice through an unconsolidated clay layer: Implications for the stability of ice on Mars and the possible diurnal variations in atmospheric water

We have studied the sublimation of ice and water vapor transport through various thicknesses of clay (<63 μm grain size). We experimentally demonstrate that both adsorption and diffusion strongly affect the transport of water, and that the processes of diffusion and adsorption can be separately quantified once the system comes to a steady state. At shallow depths of clay, water vapor transport is determined by diffusion through both the atmosphere and the clay layer, whereas at greater depth the rate of sublimation of the ice is governed only by diffusion through the clay. Using two different models, we determine the diffusion coefficient for water vapor through unconsolidated clay layer to be 1.08±0.04×¹⁰⁻⁴ and . We also determined the adsorption isotherms for the clay layer, which follow the Langmuir theory at low water vapor pressure (<100 Pa, where a monolayer of water molecules forms on the surface of the clay) and the BET theory at higher pressure (where multiple water layers form). From our analysis of both types of isotherms we determined the adsorption constants to be and c=30±10, respectively, and specific surface areas of 1.10±0.2×¹⁰⁵ and , respectively. Finally, we report a theoretical kinetic model for the simultaneous diffusion and adsorption from which we determine adsorption kinetic constants according to the Langmuir theory of and . If the martian regolith possesses diffusive properties similar to those of the unconsolidated montmorillonite soil we investigated here, it would not represent a significant barrier to the sublimation of subsurface ice. However, at the low subsurface temperatures of high latitude (180 K on average), ice could survive from the last glaciation period (about 300 to 400,000 years ago). Higher subsurface temperatures in the equatorial regions would prevent long-timescale survival of ice in the shallow subsurface. In agreement with previous work, we show that adsorption of water by a clay regolith could provide a significant reservoir of subsurface water and it might account for the purported diurnal cycle in the water content of the atmosphere.     

A global solution for the Mars static and seasonal gravity, Mars orientation, Phobos and Deimos masses, and Mars ephemeris

With the collection of six years of MGS tracking data and three years of Mars Odyssey tracking data, there has been a continual improvement in the JPL Mars gravity field determination. This includes the measurement of the seasonal changes in the gravity coefficients (e.g., , , , , , ) caused by the mass exchange between the polar ice caps and atmosphere. This paper describes the latest gravity field MGS95J to degree and order 95. The improvement comes from additional tracking data and the adoption of a more complete Mars orientation model with nutation, instead of the IAU 2000 model. Free wobble of the Mars' spin axis, i.e. polar motion, has been constrained to be less than 10 mas by looking at the temporal history of and . A strong annual signature is observed in , and this is a mixture of polar motion and ice mass redistribution. The Love number solution with a subset of Odyssey tracking data is consistent with the previous liquid outer core determination from MGS tracking data [Yoder et al., 2003. Science 300, 299-303], giving a combined solution of k₂=0.152±0.009 using MGS and Odyssey tracking data. The solutions for the masses of the Mars' moons show consistency between MGS, Odyssey, and Viking data sets; Phobos GM=(7.16±0.005)×¹⁰⁻⁴ km³/s² and Deimos GM=(0.98±0.07)×¹⁰⁻⁴ km³/s². Average MGS orbit errors, determined from differences in the overlaps of orbit solutions, have been reduced to 10-cm in the radial direction and 1.5 m along the spacecraft velocity and normal to the orbit plane. Hence, the ranging to the MGS and Odyssey spacecraft has resulted in position measurements of the Mars system center-of-mass relative to the Earth to an accuracy of one meter, greatly reducing the Mars ephemeris errors by several orders of magnitude, and providing mass estimates for Asteroids 1 Ceres, 2 Pallas, 3 Juno, 4 Vesta, and 324 Bamberga.     

Scenarios for the formation of Chasma Boreale, Mars

The martian polar caps feature large chasmata and smaller trough systems which have no counterpart in terrestrial ice sheets. Chasma Boreale cuts about 500 km into the western part of the north-polar cap, is up to 100 km wide and up to 2 km deep. One possible formation mechanism is by a temporary heat source under the ice due to tectono-thermal or volcanic activity, which melts the ice from below. It is demonstrated by model simulations that this process is feasible, a moderately increased heat flux of 0.5-1 W m⁻², sustained over at least tens of thousands of years, producing a topographic depression which resembles the real chasma. Associated meltwater discharge rates are small (), but can exceed 10 km³ a⁻¹ if a stronger heat flux of 10 W m⁻² is assumed. Local ice-flow velocities during the process of chasma formation can exceed 1 m a⁻¹ at the head and scarps of the chasma. However, if the thermal anomaly shuts down, glacial flow quickly decreases, so that the chasma can stay open for an indefinite amount of time without an ongoing, sustaining process under the climate conditions of the most recent millions of years.     

Recent volcano-ice interaction and outburst flooding in a Mars polar cap re-entrant

Formation of chasms in the polar ice caps of Mars has been attributed to meltwater outburst floods, but the cause of melting has remained uncertain. In a cap re-entrant enveloping Abalos Colles, west of Casma Boreale in the north polar cap, we have found possible evidence of recent volcano-ice interaction and outburst flooding. In this paper we demonstrate that these two mechanisms can have acted together to form or expand the Abalos re-entrant. Flat-topped ridges and circular rims protruding above the ice cap surface in the re-entrant apex may be lava ridges and volcano craters, and can have caused melting of 3.3 to 7.7×¹⁰³ km³ of ice. The surrounding cap surface appears to have subsided and the likely volume of missing ice matches the melt estimate. Outburst flooding from this area may have reached peak discharges of 0.3 to according to scour patterns in one of the re-entrant channels. This required ponding of melt water during lava eruption and catastrophic release through a sub- or englacial melt water tunnel, the collapse of which has left a chasm in the ice cap margin. The flood features are geologically recent, and volcano-ice interaction may have occurred within the last 20,000 years.     

Heavy ion escape from Mars, influence from solar wind conditions and crustal magnetic fields

We have used more than 4 years of Mars Express ion data to estimate the escape of heavy ions ( and ) from Mars. To take the limited field of view of the instrument into account, the data has been binned into spatial bins and angular bins to create average distribution functions for different positions in the near Mars space. The net escape flux for the studied low solar activity period, between May 2007 and May 2011, is 2.0 ± 0.2 × 10²⁴ s⁻¹. The escape has been calculated independently for four different quadrants in the Y_MSO − Z_MSO plane, south, dusk, north and dawn. Escape is highest from the northern and dusk quadrants, 0.6 ± 0.1 × 10²⁴ s⁻¹, and smallest from the south and dawn quadrants, 0.4 ± 0.1 × 10²⁴ s⁻¹. The flux ratio of molecular ( and ) to O⁺ ions is 0.9 ± 0.1, averaged over all quadrants. The flux difference between the north and south quadrants is statistically significant, and is presumed to be due to the presence of significant crustal magnetic fields in the southern hemisphere, reducing the outflow. The difference between the dawn and dusk quadrants is likely due to the magnetic tension associated with the nominal Parker angle spiral, which should lead to higher average magnetic tension on the dusk side. The escape increases during periods of high solar wind flux and during times when co-rotating interaction regions (CIR) affect Mars. In the latter case the increase is a factor 2.4-2.9 as compared to average conditions.     

A process to make massive ice in the martian regolith using long-term diffusion and thermal cracking

This paper describes a “simple standard” model of water transport through regolith that includes diffusive migration and phase changes driven by damped seasonal temperature waves. A hitherto unused first-order process is then added, that can produce ice densities much greater than those allowed by the initial dry porosity. Voids are produced in cooling icy regolith when tensile stresses exceed the cracking threshold . These stresses build up through an interaction of thermal contraction and elastic-plastic response. When the cracks open up after tensile failure there is purely thermal void enhancement and subsequent reduction as the regolith warms again. When the cracks are open the porosity is increased and they partially fill with ice crystals. Thus the void reduction on warming cannot go back to the original zero point and the bulk density of ice is increased with each temperature cycle. The cracking and thermal adjustment happen at scales of meters to millimeters. The large cracks can occur in pure ice and/or homogeneous icy material and the smaller cracks are produced by rock cobbles, pebbles, and grains having a much smaller coefficient of thermal expansion than ice. Thus a hierarchy of cracks and voids forms each temperature cycle and augments the ice content. The process can take the upper few meters of a pore-saturated icy soil from 28% by mass ice content to 70% in 10 Ma. This mechanism and the seasonal temperature cycle can plausibly produce massive ice deposits in the upper few meters of Mars' high-latitude regolith by diffusion and also keep the massive-ice regolith effectively porous to water vapor transport. The obliquity cycle can produce tensile stresses nearing 2 MPa down to depth so even deeper cracking could be happening.     

Comparison between the findings of Deep Impact and our experimental results on large samples of gas-laden amorphous ice

The findings of Deep Impact on the structure and composition of Tempel-1 are compared with our experimental results on large (20 cm diameter and up to 10 cm high) samples of gas-laden amorphous ice. The mechanical ∼tensile strength inferred for Tempel-1: ∼65 Pa is 30 to 60 times smaller than our experimental findings of 2-4 kPa. This means that Tempel-1 is even fluffier than our very fluffy, talcum like, ice sample. The thermal inertia: is very close to our value of 80. The density of , is close to our value of 250-300 kg m⁻³, taking into account an ice/silicate ratio of 1 in the comet, while we study pure ice. Surface morphological features, such as non-circular depressions, chaotic terrain and smooth surfaces, were observed in our experiments. The only small increase in the gas/water vapor ratio pre- and post-impact, suggest that in the area excavated by the impactor, the 135 K front did not penetrate deeper than a few meters. Altogether, the agreement between the findings of Deep Impact and our experimental results point to a loose agglomerate of ice grains (with a silicate-organic core), which was formed by a very gentle aggregation of the ice grains, without compaction.     

Extension of atmospheric dust loading to high altitudes during the 2001 Mars dust storm: MGS TES limb observations

Mars Global Surveyor (MGS) visible (solarband bolometer) and thermal infrared (IR) spectral limb observations from the Thermal Emission Spectrometer (TES) support quantitative profile retrievals for dust opacity and particle sizes during the 2001 global dust event on Mars. The current analysis considers the behavior of dust lifted to altitudes above 30 km during the course of this storm; in terms of dust vertical mixing, particle sizes, and global distribution. TES global maps of visible (solarband) limb brightness at 60 km altitude indicate a global-scale, seasonally evolving (over 190-240° solar longitudes, L_S) longitudinal corridor of vertically extended dust loading (which may be associated with a retrograde propagating, wavenumber 1 Rossby wave). Spherical radiative transfer analysis of selected limb profiles for TES visible and thermal IR radiances provide quantitative vertical profiles of dust opacity, indicating regional conditions of altitude-increasing dust mixing ratios. Observed infrared spectral dependences and visible-to-infrared opacity ratios of dust scattering over 30-60 km altitudes indicate particle sizes characteristic of lower altitudes (cross-section weighted effective radius, ), during conditions of significant dust transport to these altitudes. Conditions of reduced dust loading at 30-60 km altitudes present smaller dust particle sizes . These observations suggest rapid meridional transport at 30-80 km altitudes, with substantial longitudinal variation, of dust lifted to these altitudes over southern hemisphere atmospheric regions characterized by extraordinary (m/s) vertical advection velocities. By L_S=230° dust loading above 50 km altitudes decreased markedly at southern latitudes, with a high altitude (60-80 km) haze of fine (likely) water ice particles appearing over 10°S-40°N latitudes.     

Observations of C2H6 and C2H2 in the stratosphere of Saturn

We have performed high-resolution spectral observations at mid-infrared wavelengths of C₂H₆ (12.16 μm), and C₂H₂ (13.45 μm) on Saturn. These emission features probe the stratosphere of the planet and provide information on the hydrocarbon photochemical processes taking place in that region of the atmosphere. The observations were performed using our cryogenic echelle spectrometer Celeste, in conjunction with the McMath-Pierce 1.5-m solar telescope in November and December 1994. We used Voyager IRIS CH₄ observations (7.67 μm) to derive a temperature profile on the saturnian atmosphere for the region of the stratosphere. This profile was then used in conjunction with height-dependent volume mixing ratios of each hydrocarbon to determine global abundances for ethane and acetylene. Our ground-based measurements indicate abundances of for C₂H₆ (1.0 mbar pressure level), and for C₂H₂ (1.6 mbar pressure level). We also derived new mixing ratios from the Voyager mid-latitude IRIS observations; 8.6±0.9×¹⁰⁻⁶ for C₂H₆ (0.1-3.0 mbar pressure level), and 1.6±0.2×¹⁰⁻⁷ for C₂H₂ (2.0 mbar pressure level).     

Micrometer-sized ice particles for planetary-science experiments - I. Preparation, critical rolling friction force, and specific surface energy

Coagulation models assume a higher sticking threshold for micrometer-sized ice particles than for micrometer-sized silicate particles. However, in contrast to silicates, laboratory investigations of the collision properties of micrometer-sized ice particles (in particular, of the most abundant H₂O-ice) have not been conducted yet. Thus, we used two different experimental methods to produce micrometer-sized H₂O-ice particles, i.e. by spraying H₂O droplets into liquid nitrogen and by spraying H₂O droplets into a cold nitrogen atmosphere. The mean particle radii of the ice particles produced with these experimental methods are (1.49 ± 0.79) μm and (1.45 ± 0.65) μm. Ice aggregates composed of the micrometer-sized ice particles are highly porous (volume filling factor: ? = 0.11 ± 0.01) or rather compact (volume filling factor: ? = 0.72 ± 0.04), depending on the method of production. Furthermore, the critical rolling friction force of F_Roll,ice = (114.8 ± 23.8) × 10⁻¹⁰ N was measured for micrometer-sized ice particles, which exceeds the critical rolling friction force of micrometer-sized SiO₂ particles . This result implies that the adhesive bonding between micrometer-sized ice particles is stronger than the bonding strength between SiO₂ particles. An estimation of the specific surface energy of micrometer-sized ice particles, derived from the measured critical rolling friction forces and the surface energy of micrometer-sized SiO₂ particles, results in γ_ice = 0.190 J m⁻².     

Parametric analysis of modeled ion escape from Mars

We develop a parametric fit to the results of a detailed magnetohydrodynamic (MHD) study of the response of ion escape rates (O⁺, and ) to strongly varied solar forcing factors, as a way to efficiently extend the MHD results to different conditions. We then use this to develop a second, evolutionary model of solar forced ion escape. We treat the escape fluxes of ion species at Mars as proportional to the product of power laws of four factors - that of the EUV flux R_euv, the solar wind particle density R_ρ, its velocity (squared) R_v², and the interplanetary magnetic field pressure R_B², where forcing factors are expressed in units of the current epoch-averaged values. Our parametric model is: , where ?(i) is the escape flux of ion i. We base our study on the results of just six provided MHD model runs employing large forcing factor variations, and thus construct a successful, first-order parametric model of the MHD program. We perform a five-dimensional least squares fit of this power law model to the MHD results to derive the flux normalizations and the indices of the solar forcing factors. For O⁺, we obtain the values, 1.73 × 10²⁴ s⁻¹, 0.782, 0.251, 0.382, and 0.214, for ?₀, α, β, γ, and δ, respectively. For , the corresponding values are 1.68 × 10²⁴ s⁻¹, −0.393, 0.798, 0.967, and 0.533. For , they are 8.66 × 10²² s⁻¹, −0.427, 1.083, 1.214, and 0.690. The fit reproduces the MHD results to an average error of about 5%, suggesting that the power laws are broadly representative of the MHD model results. Our analysis of the MHD model shows that by itself an increase in R_EUV enhances O⁺ loss, but suppresses the escape of and , whereas increases in solar wind (i.e., in , and R_B², with R_euv constant) favors the escape of heavier ions more than light ions. The ratios of escaping ions detectable at Mars today can be predicted by this parametric fit as a function of the solar forcing factors. We also use the parametric model to compute escape rates over martian history. This second parametric model expresses ion escape functions of one variable (per ion), ?(i) = ?₀(i)(t/t₀)^−ξ(i). The ξ(i) are linear combinations of the epoch-averaged ion escape sensitivities, which are seen to increase with ion mass. We integrate the and oxygen ion escape rates over time, and find that in the last 3.85 Gyr, Mars would have lost about mbars of , and of water (from O⁺ and ) from ion escape.     

The E ring in the vicinity of Enceladus: I. Spatial distribution and properties of the ring particles

Saturn's diffuse E ring is the largest ring of the Solar System and extends from about (Saturn radius RS=60,330 km) to at least encompassing the icy moons Mimas, Enceladus, Tethys, Dione, and Rhea. After Cassini's insertion into her saturnian orbit in July 2004, the spacecraft performed a number of equatorial as well as steep traversals through the E ring inside the orbit of the icy moon Dione. Here, we report about dust impact data we obtained during 2 shallow and 6 steep crossings of the orbit of the dominant ring source—the ice moon Enceladus. Based on impact data of grains exceeding 0.9 μm we conclude that Enceladus feeds a torus populated by grains of at least this size along its orbit. The vertical ring structure at agrees well with a Gaussian with a full-width-half-maximum (FWHM) of ∼4200 km. We show that the FWHM at is due to three-body interactions of dust grains ejected by Enceladus' recently discovered ice volcanoes with the moon during their first orbit. We find that particles with initial speeds between 225 and 235 m s⁻¹ relative to the moon's surface dominate the vertical distribution of dust. Particles with initial velocities exceeding the moon's escape speed of 207 m s⁻¹ but slower than 225 m s⁻¹ re-collide with Enceladus and do not contribute to the ring particle population. We find the peak number density to range between 16×¹⁰⁻² m⁻³ and 21×¹⁰⁻² m⁻³ for grains larger 0.9 μm, and 2.1×¹⁰⁻² m⁻³ and 7.6×¹⁰⁻² m⁻³ for grains larger than 1.6 μm. Our data imply that the densest point is displaced outwards by at least with respect of the Enceladus orbit. This finding provides direct evidence for plume particles dragged outwards by the ambient plasma. The differential size distribution for grains >0.9 μm is described best by a power law with slopes between 4 and 5. We also obtained dust data during ring plane crossings in the vicinity of the orbits of Mimas and Tethys. The vertical distribution of grains >0.8 μm at Mimas orbit is also well described by Gaussian with a FWHM of ∼5400 km and displaced southwards by ∼1200 km with respect to the geometrical equator. The vertical distribution of ring particles in the vicinity of Tethys, however, does not match a Gaussian. We use the FWHM values obtained from the vertical crossings to establish a 2-dimensional model for the ring particle distribution which matches our observations during vertical and equatorial traversals through the E ring.     

Gravity field and interior of Rhea from Cassini data analysis

The Cassini spacecraft encountered Rhea on November 26, 2005. Analysis of the Doppler data acquired at and around closest approach yields the mass of Rhea and the quadrupole moments of its gravity field with unprecedented accuracy. We obtained which corresponds to a density of . Our results for J₂ and C₂₂ are (7.947±0.892)×¹⁰⁻⁴ and (2.3526±0.0476)×¹⁰⁻⁴, respectively. These values are consistent with hydrostatic equilibrium. From the value of C₂₂, we infer the non-dimensional moment of inertia C/MR²=0.3721±0.0036. Our models of Rhea's interior based on the gravity data favor an almost undifferentiated satellite. A discontinuity between a core and a mantle is possible but not required by the data. Models with a constant silicate mass fraction throughout the body cannot account for the determined quadrupole coefficients. The data exclude fully differentiated models in which the core would be composed of unhydrated silicates and the mantle would be composed of pure ice. If the mantle contains 10% in mass of silicates, the core extends to 630 km in radius and has a silicate mass fraction of 40%. A continuous model in which the silicates are more concentrated toward the center of the body than in the outer layers is allowed by the gravity data but excluded by thermal evolution considerations. The one model that fits the gravity data and is self-consistent when energy transport and ice melting are qualitatively considered is an “almost undifferentiated” Rhea, in which a very large uniform core is surrounded by a relatively thin ice shell containing no rock at all.     

Mars high resolution gravity fields from MRO, Mars seasonal gravity, and other dynamical parameters

With 2 years of tracking data collection from the MRO spacecraft, there is noticeable improvement in the high frequency portion of the spherical harmonic Mars gravity field. The new JPL Mars gravity fields, MRO110B and MRO110B2, show resolution near degree 90. Additional years of MGS and Mars Odyssey tracking data result in improvement for the seasonal gravity changes which compares well to global circulation models and Odyssey neutron data and Mars rotation and precession (). Once atmospheric dust is accounted for in the spacecraft solar pressure model, solutions for Mars solar tide are consistent between data sets and show slightly larger values (k₂ = 0.164 ± 0.009, after correction for atmospheric tide) compared to previous results, further constraining core models. An additional 4 years of Mars range data improves the Mars ephemeris, determines 21 asteroid masses and bounds solar mass loss (dGM_Sun/dt < 1.6 × 10⁻¹³ GM_Sun year⁻¹).     

Importance of aeolian processes in the origin of the north polar chasmata, Mars

A detailed examination of the location and orientation of sand dunes and other aeolian features within the north polar chasmata indicates that steep scarps strongly influence the direction and intensity of prevailing winds. These steep scarps are present at the heads and along the margins of the north polar chasmata. Topographic profiles of the arcuate head scarps and equator-facing wall of Chasma Boreale reveal unusually steep polar slopes ranging from ∼6°-30°. The relatively steep-sloped (∼8°), sinuous scarp at the head of two smaller chasmata, located west of Chasma Boreale, exhibits an obvious resistant cap-forming unit. Scarp retreat is occurring in places where the cap unit is actively being undercut by descending slope winds. Low-albedo surfaces lacking sand dunes or dust mantles are present at the base of the polar scarps. A ∼100-300 m deep moat, located at the base of the scarps, corresponds with these surfaces and indicates an area of active aeolian scour from descending katabatic winds. Small local dust storms observed along the equator-facing wall of Chasma Boreale imply that slope wind velocities in Chasma Boreale are sufficient to mobilize dust and sand-sized particles in the Polar Layered Deposits (PLD). Two amphitheater forms, located above the cap-forming unit of the sinuous scarp and west of Chasma Boreale, may represent an early stage of polar scarp and chasma formation. These two forms are developing within a younger section of polar layered materials. The unusually steep scarps associated with the polar chasmata have developed where resistant layers are present in the PLD, offering resistance during the headward erosion and poleward retreat of the scarps. Steep slopes that formed under these circumstances enhance the flow of down-scarp katabatic winds. On the basis of these observations, we reject the fluvial outflood hypothesis for the origin of the north polar chasmata and embrace a wind erosion model for their long-term development. In the aeolian model, off-pole katabatic winds progressively remove materials from the steep slopes below chasmata scarps, undermining resistant layers at the tops of scarps and causing retreat by headward erosion. Assuming a minimum age for the onset of formation of Chasma Boreale (10⁵ yr) results in a maximum volumetric erosion rate of . Removal of this volume of material from the equator-facing wall and head scarps of chasma would require a rate for scarp retreat of .     

CO2 production by ion irradiation of H2O ice on top of carbonaceous materials and its relevance to the Galilean satellites

In this work we report on new experiments of ion irradiation of water ice deposited on top of solid carbonaceous materials to study the production of CO₂ at the interface ice/refractory material and discuss the possibility that this mechanism accounts for the quantity of CO₂ ice detected on the surfaces of the Galilean satellites. The used experimental technique has been in situ infrared spectroscopy. We have irradiated thin films of H₂O frost on carbonaceous layers with 200 keV of He⁺ and Ar⁺, and 30 keV of He⁺ at 16 and 80 K. The used carbonaceous layers have been asphaltite, a natural bitumen, and solid organic residues obtained by irradiation of frozen benzene. In both cases the results show that CO₂ is produced very efficiently after irradiation obtaining a maximum quantity of the order of . These results are, also quantitatively similar, to those recently obtained for water ice deposited on amorphous carbon films [Mennella, V., Palumbo, M.E., Baratta, G.A., 2004. Formation of CO and CO₂ molecules by ion irradiation of water ice covered hydrogenated carbon grains. Astrophys. J. 615, 1073-1080]. Thus we suggest that, whatever is the carbonaceous residue, CO₂ will be produced efficiently by the studied process. These results have interest in the context of the surfaces of the icy Galilean satellites in which CO₂ has been detected mainly trapped in the non-ice material, not in the pure water ice. We suggest that radiolysis of mixtures of water ice and refractory carbonaceous materials is the primary formation mechanism responsible for the CO₂ formation on the surfaces of the Galilean satellites.     

Spectral imaging of martian water ice clouds and their diurnal behavior during the 1999 aphelion season (Ls = 130°)

We report high-spectral-resolution (λ/δλ = 800-2300) near-infrared mapping observations of Mars at L_s = 130° (April 1999), which were obtained by drift-scanning the cryogenic long-slit spectrometer at the KPNO 2.2-m telescope across the disk. Data were reformatted into calibrated spectral image cubes (x,y,λ) spanning 2.19 to 4.12 μm, which distinguish atmospheric CO₂ features, solar lines, and surface and aerosol features. Maps of relative band depth between 3.0 and 3.5 μm trace water ice clouds and show the diurnal evolution of features in the persistent northern summer aphelion cloud belt, which was mapped contemporaneously but at fixed local time by the Mars Global Surveyor Thermal Emission Spectrometer (MGS/TES). Cloud optical depth, particle sizes, and ice aerosol content were estimated using a two-stream, single-layer scattering model, with Mie coefficients derived from recently published ice optical constants, followed by a linear spectral deconvolution process. A comparison of data and model spectra shows evaporating nighttime clouds in the morning followed by afternoon growth of a prominent orographic cloud feature on the west flank of Elysium Mons. Cloud optical depth at 3.2 μm evolved to 0.28 ± 0.13 and ice aerosol column abundance to 0.9 ± 0.3 pr μm in the afternoon. Column abundances as large as 0.17 pr μm were retrieved in nonorographic clouds within the aphelion cloud band around midday. These clouds exhibit a modest decline in optical depth during the afternoon. Results show that ice particle radii from <2 μm to >4 μm exist in both cloud types. However, large particles dominate the spectra, consistent with recent MGS/TES emission phase function measurements of aphelion cloud aerosol properties.     

The global shape of Europa: Constraints on lateral shell thickness variations

The global shape of Europa is controlled by tidal and rotational potentials and possibly by lateral variations in ice shell thickness. We use limb profiles from four Galileo images to determine the best-fit hydrostatic shape, yielding a mean radius of 1560.8±0.3 km and a radius difference a−c of 3.0±0.9 km, consistent with previous determinations and inferences from gravity observations. Adding long-wavelength topography due to proposed lateral variations in shell thickness results in poorer fits to the limb profiles. We conclude that lateral shell thickness variations and long-wavelength isostatically supported topographic variations do not exceed 7 and 0.7 km, respectively. For the range of rheologies investigated (basal viscosities from 10¹⁴ to ) the maximum permissible (conductive) shell thickness is 35 km. The relative uniformity of Europa's shell thickness is due to either a heat flux from the silicate interior, lateral ice flow at the base of the shell, or convection within the shell.