Martian aqueous mineral deposits have been examined and characterized using data acquired during Mars Reconnaissance Orbiter's (MRO) primary science phase, including Compact Reconnaissance Imaging Spectrometer for Mars hyperspectral images covering the 0.4–3.9 μm wavelength range, coordinated with higher–spatial resolution HiRISE and Context Imager images. MRO's new high‐resolution measurements, combined with earlier data from Thermal Emission Spectrometer; Thermal Emission Imaging System; and Observatoire pour la Minéralogie, L'Eau, les Glaces et l'Activitié on Mars Express, indicate that aqueous minerals are both diverse and widespread on the Martian surface. The aqueous minerals occur in 9–10 classes of deposits characterized by distinct mineral assemblages, morphologies, and geologic settings. Phyllosilicates occur in several settings: in compositionally layered blankets hundreds of meters thick, superposed on eroded Noachian terrains; in lower layers of intracrater depositional fans; in layers with potential chlorides in sediments on intercrater plains; and as thousands of deep exposures in craters and escarpments. Carbonate‐bearing rocks form a thin unit surrounding the Isidis basin. Hydrated silica occurs with hydrated sulfates in thin stratified deposits surrounding Valles Marineris. Hydrated sulfates also occur together with crystalline ferric minerals in thick, layered deposits in Terra Meridiani and in Valles Marineris and together with kaolinite in deposits that partially infill some highland craters. In this paper we describe each of the classes of deposits, review hypotheses for their origins, identify new questions posed by existing measurements, and consider their implications for ancient habitable environments. On the basis of current data, two to five classes of Noachian‐aged deposits containing phyllosilicates and carbonates may have formed in aqueous environments with pH and water activities suitable for life.
Observations by the Mars Reconnaissance Orbiter/Compact Reconnaissance Imaging Spectrometer for Mars in the Mawrth Vallis region show several phyllosilicate species, indicating a wide range of past aqueous activity. Iron/magnesium (Fe/Mg)–smectite is observed in light-toned outcrops that probably formed via aqueous alteration of basalt of the ancient cratered terrain. This unit is overlain by rocks rich in hydrated silica, montmorillonite, and kaolinite that may have formed via subsequent leaching of Fe and Mg through extended aqueous events or a change in aqueous chemistry. A spectral feature attributed to an Fe 2+ phase is present in many locations in the Mawrth Vallis region at the transition from Fe/Mg-smectite to aluminum/silicon (Al/Si)–rich units. Fe 2+ -bearing materials in terrestrial sediments are typically associated with microorganisms or changes in pH or cations and could be explained here by hydrothermal activity. The stratigraphy of Fe/Mg-smectite overlain by a ferrous phase, hydrated silica, and then Al-phyllosilicates implies a complex aqueous history.
The Cassini visual and infrared mapping spectrometer (VIMS) investigation is a multidisciplinary study of the Saturnian system. Visual and near-infrared imaging spectroscopy and high-speed spectrophotometry are the observational techniques. The scope of the investigation includes the rings, the surfaces of the icy satellites and Titan, and the atmospheres of Saturn and Titan. In this paper, we will elucidate the major scientific and measurement goals of the investigation, the major characteristics of the Cassini VIMS instrument, the instrument calibration, and operation, and the results of the recent Cassini flybys of Venus and the Earth-Moon system.
Spectra from Cassini's Visual and Infrared Mapping Spectrometer reveal that the horizontal structure, height, and optical depth of Titan's clouds are highly dynamic. Vigorous cloud centers are seen to rise from the middle to the upper troposphere within 30 minutes and dissipate within the next hour. Their development indicates that Titan's clouds evolve convectively; dissipate through rain; and, over the next several hours, waft downwind to achieve their great longitude extents. These and other characteristics suggest that temperate clouds originate from circulation-induced convergence, in addition to a forcing at the surface associated with Saturn's tides, geology, and/or surface composition.
Titan is the only satellite in our Solar System with a dense atmosphere. The surface pressure is 1.5 bar (ref. 1) and, similar to the Earth, N2 is the main component of the atmosphere. Methane is the second most important component, but it is photodissociated on a timescale of 10(7) years (ref. 3). This short timescale has led to the suggestion that Titan may possess a surface or subsurface reservoir of hydrocarbons to replenish the atmosphere. Here we report near-infrared images of Titan obtained on 26 October 2004 by the Cassini spacecraft. The images show that a widespread methane ocean does not exist; subtle albedo variations instead suggest topographical variations, as would be expected for a more solid (perhaps icy) surface. We also find a circular structure approximately 30 km in diameter that does not resemble any features seen on other icy satellites. We propose that the structure is a dome formed by upwelling icy plumes that release methane into Titan's atmosphere.
[1] Visible-near infrared reflectance spectra acquired by the Mars Express Observatoire pour la Minéralogie, l'Eau, les Glaces et l'Activité (OMEGA) spectrometer are used to estimate the absolute water content within the uppermost fraction of the Martian regolith. This upper surface layer represents the boundary between the regolith and atmosphere; thus the amount of water stored in these two reservoirs and at this boundary is expected to vary spatially and temporally with changing equilibrium conditions. We have applied models derived from laboratory experiments relating the strength of the 3 mm hydration feature to absolute water content to OMEGA spectra acquired during the first $1200 orbits of the Mars Express mission to estimate the H 2 O content of the Martian surface. Three methods were used to examine the strength of the 3 mm absorption: integrated band depth, apparent absorbance, and the effective single-particle absorptionthickness (ESPAT) parameter. Integrated band depth and apparent absorbance values are correlated to albedo when derived from reflectance spectra, implying that bright regions are more hydrated than dark regions. The ESPAT parameter, however, relies on single scattering albedo instead of reflectance and is capable of estimating absolute water content within ±1 wt.% H 2 O for a wide range of albedo values, compositions, and particle sizes. Applying this model to the OMEGA data reveals that bright and dark regions commonly have similar water contents in equatorial regions and the largest spatial variations in H 2 O occur as a function of latitude. Equatorial regions exhibit water contents in the range of $2-5 wt.%, whereas latitudes higher than $45°N are characterized by a continuous increase in H 2 O with latitude from $5-15 wt.%. Phyllosilicate and sulfate bearing terrains are more hydrated than average bright and dark regions and their locations are in widely separated areas of Noachian-aged material, suggesting chemical alteration by water-rock interaction may have been spatially extensive in the early history of Mars.
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