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The spectral analysis of a planetary surface is fundamental for a deeper understanding of the mineralogy and composition. In particular, the determination of spectral units is a reliable method to infer the physical and compositional properties of a surface by processing several spectral parameters simultaneously, instead of the more traditional approach of interpreting each single parameter separately. To define the spectal units, we first compute the most relevant spectral parameters, based on a preliminary detailed analysis of the spectral properties of a surface. This method could be used for different bodies and is described in [1].For this work, we selected the Apollo Basin area within South Polar Aitken [2,3], the largest and deepest impact basin on the Moon. We analyzed the M3/Chandrayaan-1 data [4] after performing the most up-to-date calibration, thermal removal and photometric correction [5,6]. Lunar spectra are characterized by two strong pyroxenes absorption bands at 1 and 2 µm. In this regard, we decided to define the Apollo Basin spectral units by using the two pyroxenes band depths, the reflectance at 540 nm (standard visible wavelength), and the spectral slope of the 1 µm (see [7]). In Apollo Basin, we found 12 different spectral units. Among these units, the most peculiar is the one linked to the basaltic smooth plains within the floor of the crater. This unit is characterized by low reflectance, deep band depths and a strongly positive spectral slopes (more red surfaces). Subsequently, an analysis of absorption band center at 1 and 2 µm and a comparison with RELAB synthetic pyroxenes [8] revealed a composition compatible with material dominated by strong pyroxene absorptions, e.g. clinopiroxenes, such as pigeonite or augite, with Low Ca and Mg, and relatively high Fe (Fs: 34-75; En: 6-23; Wo: 10-27). The rest of the units show a similar mineralogy to the orthopyroxenes, with intermediate amount of Fe and Mg.This work allows for a detailed understanding of the mineralogy of Apollo Basin, but also lays the groundwork to search for a link between spectral, and morpho-stratigraphic units [9] to reach out highly informative geological maps of the Moon. This innovative approach is one of the main goals of the H2020 no. 776276-PLANMAP project [10].Acknowledgments: This work is funded by the European Union’s Horizon 2020 research grant agreement No 776276- PLANMAP.References: [1] Zambon et al., 2020 LPSC. [2] Ohtake, M. et al., 2014, GRL. [3] Moriarty, D.P. et al., 2018, JGR. [4] Pieters et al., 2009, Current Science. [5] PLANMAP D4.3- Spectral Indices and RGB maps. [6] Besse, S. et al., 2012, Icarus. [7] PLANMAP D4.3- Spectral Indices and RGB maps. [8] http://planetary.brown.edu/relabdocs/synth_pyx/pyroxenes.html. [9] Ivanov, M.A., 2018, JGR. [10] https://www.planmap.eu/.

South Polar Aitken (SPA) is the largest and deepest basin on the Moon [1, 2], and several regions within this widespread basin have been considered as possible landing sites for present and future missions. Some of these sites are also targets of the H2020 n° 776276- PLANMAP project [3]. In this work, we focus on specific areas within SPA, Apollo basin, Lebinitz and Von Karman craters regions (see fig. 1). For our analysis, we considered the M3/Chandrayaan-1 data [4], after the application of the last calibration, thermal removal and photometric correction [5, 6]. SPA is rich in high Ca-pyroxenes, except for some central crater peaks, where low Ca-pyroxenes dominate, and no extensive olivine-rich areas are observed [1, 7, 8]. Based on the spectral characteristics on the M3 spectra, we define specific spectral parameters, such as band centers, depths, spectral slopes, including selected RGB color composite combinations [5]. These indices are the starting point for retrieving the spectral units of the regions considered. The final aim of this work is to find a link between spectral, and morpho-stratigraphic units [9] to reach out highly informative geological maps of the Moon. This innovative approach holds the efforts of different research units and is one of the main goals of the PLANMAP project. References: [1] Ohtake, M. et al., 2014, GRL. [2] Moriarty, D.P. et al., 2018, JGR. [3] https://www.planmap.eu/, [4] Pieters et al., 2009, Current Science., [5] PLANMAP D4.3- Spectral Indices and RGB maps. [6] Besse, S. et al., 2012, Icarus. [7] Dhingra, D., 2018, Geosciences. [8] Melosh, H. J., et al., 2017, Geology. [9] Ivanov, M.A., 2018, JGR. Acknowledgments: This work is funded by the European Union’s Horizon 2020 research grant agreement No 776276- PLANMAP. Figure 1: Apollo basin (top-left) and Leibnitz and Von Karman craters (bottom-left) reflectance map at 540 nm. Apollo basin (top-right) and Leibnitz and Von Karman craters (bottom-right) Clementine-like mosaic derived by M3 data (bottom-right).

International audience; The Taurus-Littrow valley on the Moon was the location of intensive geologic fieldwork during three days in December 1972. In situ activities at sampling stations were systematically documented by the astronauts using a series of overlapping images taken with their Hasselblad cameras. We investigated how this Apollo image archive can be used to perform 3-D reconstructions of several boulders of interest using close-range photogrammetry. We specifically focused on seven different boulders located at Stations 2, 6, and 7, at the foot of South and North Massifs, respectively. These boulders represent samples from highland materials, which rolled down the slopes of the surrounding hills. We used the Agisoft Metashape software to compute 3-D reconstructions of these boulders, using 173 scanned images as input. We then used either a web-based platform or a game engine to render the models in virtual reality. This allowed the users to walk around the boulders and to investigate in detail their morphology, fractures, vesicles, color variations, and sampling spots, as if standing directly in front of them with the astronauts. This work suggests that many features can be reconstructed in other sites of the Apollo missions, so as other robotic landing sites. Virtual reality techniques coupled to photogrammetry is thus opening a new era of exploration, both for past and future landing sites.

International audience; So far, only two ice species have been identified by Rosetta/VIRTIS-M [1] on the surface of 67P/Churyumov-Gerasimenko during the pre-perihelion time: crystalline water and carbon dioxide ice. Water ice has been spectroscopically identified in three distinct modalities: 1) On the active areas of Hapi region where water ice changes its abundance with local time and illumination conditions, condensing during the night hours and sublimating during daytime [2]; 2) On recent debris fields collapsed from two elevated structures in the Imhotep region where more fresh and pristine material is exposed [3]; 3) On eight bright areas located in Khonsu, Imhotep, Anhur, Atum and Khepry regions [4] where single or multiple grouped icy patches with sizes ranging between few meters to about 60 m are observed. Carbon dioxide ice has been detected only in a 60-80 m area in Anhur region while it was exiting from a four year-long winter-night season [5]. This ice deposit underwent a rapid sublimation, disappearing in about one month after its initial detection. While water and carbon dioxide ice appear always mixed with the ubiquitous dark material [6,7], there are no evidences of the presence of water and carbon dioxide ices mixed together in the same area. If observed, ices always account for very small fraction (few percent) with respect to the dark material. Moreover, the surface ice deposits are preferentially located on the large lobe and the neck while they are absent on the small lobe. Apart from these differences in the spatial distribution of ices on the surface, a large variability is observed the mixing modalities and in the grain size distributions, as retrieved from spectral modeling [8]: 1) very small Îm-sized water ice grains in intimate mixing with the dark terrain are detected on Hapi active regions [2]; 2) two monodispersed distributions with maxima at 56 Îm and at 2 mm, corresponding to the intimate and areal mixing classes, are observedon the Imhotep debris fields ice grains [3]; 3) different combinations of water ice and dark terrain in intimate mixing with small grains (tens of microns) or in areal mixing with large grains (mm- sized) are seen on the eight bright areas discussed in [4]; 4) the CO2 ice in the Anhur region appears grouped in areal patches made of 50 mum sized grains [5]. While the spectroscopic identification of water and carbon dioxide ices is made by means of diagnostic infrared absorption features, their presence cause significant effects also at visible wavelengths, including the increase of the albedo and the reduction of the spectral slope which results in a more blue color [9,10]. In summary, thermodynamic conditions prevailing on the 67P/CG nucleus surface allow the presence of only H2O and CO2 ices. Similar properties are probably common among other Jupiter family comets.

Acetylene is thought to be abundant on Titan according to most photochemical models. While detected in the atmosphere, its likely presence at the surface still lacks physical evidence. It is thought that solid acetylene could be a major component of Titan’s lakes shorelines and dry lakebed, detected as the 5 μm-bright deposits with the Cassini/VIMS instrument. Acetylene could also be present under its liquid form as dissolved solids in Titan’s methane–ethane lakes, as emphasized by thermodynamics studies. This paper is devoted to the near-infrared spectroscopy study of acetylene under solid and liquid phases between 1 and 2.2 μm, synthesized in a Titan simulation chamber that is able to reproduce extreme temperature conditions. From experiments, we observed a ∼10% albedo increase between liquid acetylene at 193–188 K and solid acetylene at 93 K. Using the NIR spectroscopy technique we successfully calculated the reflectivity ratio of solid/liquid acetylene as 1.13. The second difference we observed between liquid and solid acetylene is a shift in the major absorption band detected at 1.54 μm, the shift of ∼0.01 μm occurring toward higher wavelength. In order to assess the detectability of acetylene on Titan using the Cassini/VIMS instrument, we adapted our spectra to the VIMS spectral resolution. The spectral band at 1.55 μm and a negative slope at 2.0 μm falls in the Cassini/VIMS atmospheric windows over several VIMS infrared spectels, thus Cassini/VIMS should be able to detect acetylene.

International audience; The surface composition of Titan is of great importance for understanding both the internal evolution of Titan and its atmosphere. The Visual and Infrared Mapping Spectrometer (VIMS) investigation on Cassini is observing Titan from 0.35 to 5.11 µm with spatial resolution down to a few kilometers during each flyby of the spacecraft as it orbits Saturn. Our search for spectral diversity using seven methane transmission windows in the near infrared suggests that spectrally distinct units exist on the surface of Titan and that most of the surface can be modeled using only a few distinct spectral units: water frost, CO 2 frost, atmospheric scattering, and an unknown material bright at 2 µm. A dark, spectrally neutral material is also implied. Use of an atmospheric scattering component with spectral mixing analysis may provide a method for partially removing atmospheric effects. In some locations, atmospheric scattering accounts for the majority of the signal. There are also small regions with unusual spectra that may be due to low signal and high noise and/or may be exotic materials of interest. Further, we searched within the methane windows for spectral features associated with Titan's surface. Only the 5-µm and, to a lesser extent, the 2-µm window provide a reasonable opportunity for this, as the shorter-wavelength windows are too narrow and the 2.8-µm window is cluttered with an unknown atmospheric constituent. We find evidence for only one spectral feature: near 4.92 µm for the 5-µm bright Tui Regio region. CO 2 frost with grains smaller than about 10 µm is the best candidate we have found so far to explain this absorption as well as the feature's spectral contrast between the 2.7-and the 2.8-µm atmosphere subwindows. This suggested CO 2 identification is supported by the presence of an endmember in the spectral mixture analysis that is consistent with CO 2 frost with large grain sizes. We find no other absorption features that are statistically significant, including those reported earlier by others. These results are consistent with but greatly extend our early analysis that treated only the T a data set [McCord, T.B., et al., 2006a. Planet. Space Sci. 54, 1524-1539]. In the spectral feature search process, we explored in detail the noise characteristics of the VIMS data within the 5-µm window, which has generally very low signal (4-20 DN), due to the measurement conditions and low illumination levels. We find noise of nearly Gaussian statistics except for some erratic darks and noise spikes, and the data set seems generally well behaved. We present examples of our attempt to improve on the standard VIMS pipeline data calibration.

The surface composition of the northern regions of Vesta, observed by the Dawn spacecraft, offers the possibility to test several hypotheses related to impact-related processes. We used mostly imaging spectrometry in the visible and near infrared to assess the distribution of mafic lithologies, hydrated components and albedo properties, and use the link with howardite, eucrite and diogenite meteorites (HEDs) to investigate the origin of those materials. We established that Rheasilvia ejecta reached part of the northern regions, and have a diogenitic-rich composition characteristic of the lower crust. Investigations of the antipodes of the two major impact basins (Rheasilvia and Veneneia) did not reveal any correlation between geographic location, geological features and the surface composition. The northern wall of Mamilia crater, which is one of the freshest craters above 22°N, contains relatively pure eucritic-rich, diogenitic-rich and dark, hydrated materials, which are representative of the rest of the northern regions (and most of Vesta), with the exception of an olivine-like component found in Bellicia crater by Ammannito et al. (Ammannito, E. et al. [2013a]. Nature 504(7478), 122–125). We determined that similar types of materials are found in various proportions over a large region, including Bellicia, Arruntia and Pomponia craters, and their origin does not seem to be related to Rheasilvia ejecta. These materials are hydrated, which could indicate an exogenous origin, and not as dark as expected for carbonaceous chondrites, which likely compose the majority of dark hydrated materials on Vesta. Spectral mixture analysis reveals that mixtures of pyroxenes (hypersthene, pigeonite and diopside) could offer an alternative interpretation to olivine in this area.

International audience; The surface material of comet nucleus 67P/Churyumov-Gerasimenko contains components with a broad absorption band between 2900 and 3600 nm [1] that was revealed by the Visible and Infrared Thermal Imaging Spectrometer (VIRTIS, [2]) on the Rosetta spacecraft. Absorptions in this range of wavelengths are likely due to C-H and/or O-H chemical groups in various compounds [1, 3, 4]. In this study we investigate the distribution of this absorber across the surface of the comet nucleus, and apparent temporal variability on areas where the absorption band is the strongest.

International audience; Water-ice-poor, 5-lm-bright material on Saturn's moon Titan has previously been geomorphologically identified as evaporitic. Here we present a global distribution of the occurrences of the 5-lm-bright spectral unit, identified with Cassini's Visual Infrared Mapping Spectrometer (VIMS) and examined with RADAR when possible. We explore the possibility that each of these occurrences are evaporite deposits. The 5-lm-bright material covers 1% of Titan's surface and is not limited to the poles (the only regions with extensive, long-lived surface liquid). We find the greatest areal concentration to be in the equatorial basins Tui Regio and Hotei Regio. Our interpretations, based on the correlation between 5-lm-bright material and lakebeds, imply that there was enough liquid present at some time to create the observed 5-lm-bright material. We address the climate implications surrounding a lack of evaporitic material at the south polar basins: if the south pole basins were filled at some point in the past, then where is the evaporite?

We investigate the sulfate and iron oxide deposits in Ophir Chasma, Mars, based on short-wave infrared data from the Compact Reconnaissance Imaging Spectrometer for Mars – CRISM and from the Observatoire pour la Mineralogie, l’Eau, les Glaces et l’Activite – OMEGA. Sulfates are detected mainly in two locations. In the valley between Ophir Mensa and the southern wall of Ophir Chasma, kieserite is found both within the slope of Ophir Mensa, and superposed on the basaltic wall of the chasm. Here, kieserite is unconformably overlain by polyhydrated sulfate deposits and iron oxides. Locally, jarosite and unidentified phases with absorptions at 2.21 μm or 2.23 μm are detected, which could be mixtures of jarosite and amorphous silica or other poorly crystalline phases. The second large sulfate-rich outcrop is found on the floor of the central valley. Although the same minerals are found here, polyhydrated sulfates, kieserite, iron oxides, and locally a possibly jarosite-bearing phase, this deposit is very distinct. It is not layered, almost horizontal, and located at a much lower elevation of below −4250 m. Kieserite superposes polyhydrated sulfate-rich deposits, and iron oxides form lags. The facies of sulfate formation remains unclear, and could be different for the two locations. A formation in a lake, playa or under a glacier is consistent with the mineralogy of the central valley and its flat, low-lying topography. This is not conceivable for the kieserite deposits observed south of Ophir Mensa. These deposits are observed over several thousands of meters of elevation, which would require a standing body of water several thousands of meters deep. This would have lead to much more pervasive sulfate deposits than observed. These deposits are therefore more consistent with evaporation of groundwater infiltrating into previously sulfate-free light-toned deposits. The overlying polyhydrated sulfates and other mineral phases are observed in outcrops on ridges along the slopes of the southern chasm wall, which are too exposed to be reached by groundwater. Here, a water supply from the atmosphere by rain, snow, fog or frost is more conceivable.

International audience; Much of our knowledge of planetary surface composition is derived from remote sensing over the ultraviolet through infrared wavelength ranges. Telescopic observations and, in the past few decades, spacecraft mission observations have led to the discovery of many surface materials, from rock-forming minerals to water ice to exotic volatiles and organic compounds. Identifying surface materials and mapping their distributions allows us to constrain interior processes such as cryovolcanism and aqueous geochemistry. The recent progress in understanding of icy satellite surface composition has been aided by the evolving capabilities of spacecraft missions, advances in detector technology, and laboratory studies of candidate surface compounds. Pioneers 10 and 11, Voyagers I and II, Galileo, Cassini and the New Horizons mission have all made significant contributions. Dalton (Space Sci. Rev., 2010, this issue) summarizes the major constituents found or inferred to exist on the surfaces of the icy satellites (cf. Table 1 from Dalton, Space Sci. Rev., 2010, this issue), and the spectral coverage and resolution of many of the spacecraft instruments that have revolutionized our understanding (cf. Table 2 from Dalton, Space Sci. Rev., 2010, this issue). While much has been gained from these missions, telescopic observations also continue to provide important constraints on surface compositions, especially for those bodies that have not yet been visited by spacecraft, such as Kuiper Belt Objects (KBOs), trans-Neptunian Objects (TNOs), Centaurs, the classical planet Pluto and its moon, Charon. In this chapter, we will discuss the major satellites of the outer solar system, the materials believed to make up their surfaces, and the history of some of these discoveries. Formation scenarios and subsequent evolution will be described, with particular attention to the processes that drive surface chemistry and exchange with interiors. Major similarities and differences between the satellites are discussed, with an eye toward elucidating processes operating throughout the outer solar system. Finally we discuss the outermost satellites and other bodies, and summarize knowledge of their composition. Much of this review is likely to change in the near future with ongoing and planned outer planet missions, adding to the sense of excitement and discovery associated with our exploration of our planetary neighborhood.

International audience; This is a data paper designed to facilitate the use of and comparisons to Cassini/visual and infrared mapping spectrometer (VIMS) spectral mapping data of Saturn's moon Titan. We present thumbnail orthographic projections of flyby mosaics from each Titan encounter during the Cassini prime mission, 2004 July 1 through 2008 June 30. For each flyby we also describe the encounter geometry, and we discuss the studies that have previously been published using the VIMS dataset. The resulting compliation of metadata provides a complementary big-picture overview of the VIMS data in the public archive, and should be a useful reference for future Titan studies.

The surface of Titan has been revealed by Cassini observations in the infrared and radar wavelength ranges as well as locally by the Huygens lander instruments. Sand seas, recently discovered lakes, distinct landscapes and dendritic erosion patterns indicate dynamic surface processes. This study focus on erosional and depositional features that can be used to constrain the amount of liquids involved in the erosional process as well as on the compositional characteristics of depositional areas. Fluvial erosion channels on Titan as identified at the Huygens landing site and in RADAR and Visible and Infrared Mapping Spectrometer (VIMS) observations have been compared to analogous channel widths on Earth yielding average discharges of up to 1600 m3/s for short recurrence intervals that are sufficient to move centimeter-sized sediment and significantly higher discharges for long intervals. With respect to the associated drainage areas, this roughly translates to 1–150 cm/day runoff production rates with 10 years recurrence intervals and by assuming precipitation this implies 0.6–60 mm/h rainfall rates. Thus the observed surface erosion fits with the methane convective storm models as well as with the rates needed to transport sediment. During Cassini's T20 fly-by, the VIMS observed an extremely eroded area at 30° W, 7° S with resolutions of up to 500 m/pixel that extends over thousands of square kilometers. The spectral characteristics of this area change systematically, reflecting continuous compositional and/or particle size variations indicative of transported sediment settling out while flow capacities cease. To account for the estimated runoff production and widespread alluvial deposits of fine-grained material, release of area-dependent large fluid volumes are required. Only frequent storms with heavy rainfall or cryovolcanic induced melting can explain these erosional features.

The surface of Enceladus consists almost completely of water ice. As the band depths of water ice absorptions are sensitive to the size of particles, absorptions can be used to map variations of icy particles across the surface. The Visual and Infrared Mapping Spectrometer (VIMS) observed Enceladus with a high spatial resolution during three Cassini flybys in 2005 (orbits EN 003, EN 004 and EN 011). Based on these data we measured the band depths of water ice absorptions at 1.04, 1.25, 1.5, and 2 μm. These band depths were compared to water ice models that represent theoretically calculated reflectance spectra for a range of particle diameters between 2 μm and 1 mm. The agreement between the experimental (VIMS) and model values supports the assumption that pure water ice characterizes the surface of Enceladus and therefore that variations in band depth correspond to variations in water ice particle diameters. Our measurements show that the particle diameter of water ice increases toward younger tectonically altered surface units with the largest particles exposed in relatively “fresh” surface material. The smallest particles were generally found in old densely cratered terrains. The largest particles (∼0.2 mm) are concentrated in the so called “tiger stripes” at the south pole. In general, the particle diameters are strongly correlated with geologic features and surface ages, indicating a stratigraphic evolution of the surface that is caused by cryovolcanic resurfacing and impact gardening.

International audience; The mineralogical composition of the Martian surface is constrained by analyzing the data of the OMEGA visible and near infrared imaging spectrometer onboard Mars Express. Ferric signatures had previously been reported in Valles Marineris, Margaritifer Terra, and Terra Meridiani. Here we use three independent data reduction methods (Spectral Angle Mapper, a modified Spectral Mixture Analysis and Modified Gaussian Model) to detect and map ferric oxides in East Candor Chasma, a part of Valles Marineris. Ferric oxides in East Candor Chasma are concentrated in scattered formations. MOLA altimetry indicates that the ferric oxides are preferentially located in topographic lows. THEMIS, HRSC and MOC images show that the ferric oxide spectral signatures are systematically correlated with superficial deposits of low albedo, located at the foot of, or resting on Interior Layered Deposits (ILDs). This spatial distribution suggests that ferric oxides are genetically linked to ILDs. Gravity and wind‐driven remobilization of ferric oxides previously formed in the ILDs can explain their accumulation around the ILDs.