TP14
Session assets
Introduction: Jezero Crater lies in one of the most extensive carbonate-rich regions on Mars [1, 2], and diverse carbonate-bearing units were detected from orbit within the crater [3-6]. In particular, a unit along the Western rim of the crater stands out with especially strong carbonate signatures in CRISM data. Several origins have been considered for these “Marginal Carbonates”, including a lacustrine beach-deposit, altered igneous or pyroclastic deposits, aeolian or fluvial-deltaic origins [4]. Since September 2023 (Martian day or “sol” 910), the Perseverance rover has been exploring this so-called Margin Unit (Fig. 1A) [7].
Method: We use data collected with the SuperCam instrument [8, 9] during the Margin Campaign (up to sol 1124), with three complementary spectroscopy techniques for chemical (laser-induced breakdown spectroscopy, LIBS) and mineralogical characterization (Raman and infrared reflectance spectroscopy, IRS) of rocks along the rover traverse. Additionally, the Remote Micro-Imager (RMI) provides high-resolution color images to contextualize the spectroscopic analyses.
In particular, we derive abundances for major elements (MOC) from LIBS [10], and identify specific minerals, in particular olivine and pyroxene, based on stoichiometric analyses [11]. Carbonates (above ~50 vol% [12]) can be identified based on a combination of MOC and carbon signal characterization [12, 13]. With Raman, the position of the carbonate mode can be used to derive the composition of Fe-Mg carbonates (Mg# defined as Mg/(Fe+Mg)) [14].
Results:
LIBS – The Margin Unit shows the highest concentration of both carbonate-bearing and high silica points along the entire traverse (49% of carbonate detections and 73% of points with SiO2 > 65 wt.% for 22% of LIBS analyses). Besides these, rocks in the Margin Unit are generally of mafic composition, including multiple detections of olivine and pyroxene (Fig. 1 C-F, [11]). Although all the carbonates are Ca-poor and Fe-Mg rich, their composition is variable. In particular, before sol 1027, carbonates covered a large range of composition (Mg# varying between ~0.3 and 0.8); since sol 1027, the Mg# of identified carbonate-bearing points are clustered in the 0.6-0.8 range (Fig. 1D, F). Since these compositional changes are correlated with changes in rock morphologies and textures [7, 15, 16], we defined two sub-units: the Eastern Margin (EM) before sol 1027 and the Western Margin (WM) afterwards (Fig. 1B).
IRS – With IRS, the Margin Unit is characterized by three main signatures: i) the 1.9 μm hydration band, comparable to what is observed in the vast majority of targets in Jezero Crater [17, 18]; ii) a 2.2 μm band, more frequent and deeper than in previous units and attributed to Si-OH based on the correlation with LIBS data; iii) absorption bands at 2.3 and 2.5 μm corresponding to carbonates, with a possible contribution of clays. Most spectra also show a positive slope between 1.3 and 1.8 μm, attributed to Fe2+ in either olivine or carbonate. Additionally, a small band at 2.39 μm indicates the presence of Fe-phyllosilicates in some targets, but their precise characterization is complicated by the strong carbonate signatures.
Raman – For the first time in the mission, carbonates were detected not only in all three abraded patches (Fig. 1B) but also on six different targets with natural surfaces in the Western Margin. Some significant variability in Mg# is observed (median values: EM=~0.45; WM=~0.6).
With Raman, olivine was also detected in the first abrasion patch in the EM. No high-silica phase was detected with SuperCam Raman, which however is challenging due to an instrumental artifact.
Summary & Discussion: Several lines of evidence indicate that the Eastern and Western Margins may be distinct. Observations of layering and clastic textures show that the EM is likely of sedimentary origin. The WM lacks clear structure to conclude confidently [7, 15, 16].
The Margin Unit is uniquely enriched in both Fe-Mg carbonates and a silica-rich phase (LIBS, Raman and IRS). The silica is at least hydroxylated based on the IRS, but its hydration state is not well constrained. Additionally, points with mafic compositions in LIBS, including olivine and pyroxene grains, are also omnipresent in this unit. Beyond these general observations, we note some variability within the Margin Unit, and in particular between the Eastern and Western parts. Based on LIBS data, the WM presents a more restrained range of carbonate compositions, with only Mg-rich carbonates, covering the same range of composition as the primary minerals identified in this sub-unit (Fig. 1F). Raman data confirm the higher Mg# in carbonates in the WM. In the EM, the range of carbonate composition is wider, and wider than the range of primary mineral compositions (Fig. 1D).
The combination of mafic minerals, high-silica points and carbonates is consistent with altered mafic/ ultramafic rocks, possibly analog to carbonated peridotite [e.g. 19, 20]. The SuperCam data do not enable to confidently conclude regarding the presence of serpentine or talc in these rocks, although some phyllosilicates are detected with IRS. Nonetheless, the match between the compositions of carbonates and primary minerals in the WM is consistent with in situ alteration of olivine-rich lithology with CO2-rich fluids in a closed-system. However, the original process of emplacement of this mafic material is still poorly constrained. In the EM, the more diverse compositions – in particular of carbonates – and layering could indicate additional reworking of this altered-mafic material, possibly including Ca-enriched fluids.
References: [1] Ehlmann, Mustard et al. (2008) Nature geosc., [2] Mandon et al. (2020), [3] Goudge et al. (2017), [4] Horgan et al. (2020), [5] Tarnas et al. (2021), [6] Zastrow & Glotch (2021), [7] Horgan et al., (2024) Mars X, [8] Maurice et al. (2021) SSR. [9] Wiens al. (2021) SSR. [10] Anderson et al. (2022) SAB. [11] Udry et al., (2024) Mars X, [12] Beck et al. (2024) Icarus. [13] Clavé et al. (2023) JGR. [14] Beck & Beyssac et al., sub. [15] Jones et al., (2024) Mars X,[16] Ravanis et al., (2024) Mars X, [17] Mandon et al. (2023) JGR, [18] Dehouck et al. (2023) LPSC, [19] Beinlich et al. (2010) [20] Johnson et al. (2019)
How to cite: Clavé, E., Beck, P., Beyssac, O., Forni, O., Schröder, S., Mangold, N., Royer, C., Mandon, L., Dehouck, E., Le Mouélic, S., Quantin-Nataf, C., Udry, A., Bedford, C., Rammelkamp, K., Clegg, S., Gasnault, O., Wiens, R., and Cousin, A. and the SuperCam Team: New constraints on the “Marginal Carbonates” from in situ observations with SuperCam, Mars2020 , Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-255, https://doi.org/10.5194/epsc2024-255, 2024.
Introduction:
The study of clay-rich deposits is key to understanding past water activity on Mars and its early habitability potential. Strong phyllosilicate signatures have been identified over the Mawrth Vallis plateau based on the OMEGA instrument [1]. Similarly, large clay-rich deposits have been detected in Oxia Planum [2], the landing site of the ExoMars Rosalind Franklin rover mission. The rover aims to investigate the Martian surface and subsurface with its two-meter-deep drill, searching for potential traces of past life preserved in these phyllosilicate-bearing units [3].
Aqueous alteration minerals from the selected landing site exhibit spectral features consistent with Fe/Mg-rich phyllosilicates, best fitted with vermiculite or saponite due to the position & shape of the 2.3µm absorption band, and a Fe2+ oxidation upward slope from 1µm to ~1.7μm [2]. These large clay units date back to the Noachian period [4]. 300km northeast, clays found in Mawrth Vallis are consistent with montmorillonite (Al rich) and nontronite (Fe3+/Al rich) smectites, as indicated by spectral absorptions at respectively ~2.2µm and ~2.3µm, as well as their overall shape [5]. These units are also dated to the Noachian period [6].
The proximity of both sites and their similar position straddling the crustal dichotomy calls for the investigation of the relationship between these two significant phyllosilicate-bearing regions.
Method:
Datasets: Given the expansive scope of our study area spanning from Oxia Planum to Mawrth Vallis, we used the OMEGA hyperspectral dataset (~300 to 4000 m/pix) [7].
Tools: We developed custom Python scripts to correct OMEGA cubes, map spectral criteria, and project raster data. Extraction of regions of interest (ROIs) and spectral analysis was done in QGIS using the EnMAP-Box plugin [8].
Cube corrections: Six correction steps have been implemented: (1) an atmospheric absorption band correction (“volcano-scan” method) [9]; (2) a Mars surface thermal contribution correction (gray body subtraction) [10]; (3) corrupted pixels and spectels are removed from the cube; (4) each spectrum is fitted and then divided by a linear regression continuum with tie points centered around 1.75µm and 2.14µm, where few clay minerals absorb [11]; (5) each spectrum is normalized by its maximum value to ensure a cube with homogeneous reflectance values and mitigate significant photometric effects resulting from changes in topography; (6) the local neutral mineralogy is removed using a “clean mean” method: the median spectrum of a 50-pixel-squared region – where pixels positive to the 2.3µm criteria have been masked, hence the “clean” median – is calculated and subtracted from each spectrum. This step improves the detectability of phyllosilicate absorption bands, as their local “neutral” component is removed.
Spectral criteria maps: We mapped the following criteria: 1.9µm band depth of most hydrated minerals, 2.2µm drop of Al-smectite, 2.3µm drop of Fe-/Mg-phyllosilicates [11] (Figure 1), and HCP & LCP pyroxenes [12] (green in Figure 2).
ROIs: We converted clusters of contiguous pixels where the detection criterion was higher than a predefined threshold to shapefiles using QGIS functions. Spectral analysis over a ROI was possible after making adaptations to the EnMAP-Box Plugin [8], which were contributed to the project: by selecting one (or multiple) polygons, we compute their median spectrum (Figure 2).
Figure 1: 2.3µm-drop spectral criteria (1% threshold). Basemap: THEMIS-day.
Results:
Our 2.3µm drop criterion map aligns with existing phyllosilicate maps of Mawrth Vallis and Oxia Planum (Figure 1) [13]. Spectral analysis reveals two distinct clay types, as introduced earlier: the “Oxia Planum”-like (red in Figure 2) and the “Mawrth Vallis”-like (blue in Figure 2).
Oxia Planum clays in Mawrth Vallis and vice-versa: We noticed two areas (ROI n°1 and n°4) exhibiting different characteristics than the region they belong to. Spectrum n°1 (in Mawrth Vallis) shows Oxia Planum-like features: Fe2+ upward slope up from 1µm to 1.6µm, and the position & shape of the 2.3µm absorption. Reflectance bumps at 1.6µm (instead of 1.7µm) and 2.6µm (instead of a plateau) can be explained by the presence of pyroxenes in this ROI, which lies above a green zone, as seen in Figure 2. Spectrum n°4 (in Oxia Planum) shows the same shape & absorptions as the Mawrth-Vallis one (n°3). Spectra n°2 and n°3 show less spikes, as signal-to-noise ratio increases with the square root of the pixel count.
Lava flows covering clays? Clay detections seem to be anticorrelated with dust (dark shades in Figure 2) and, often, pyroxenes (or with traces of it in their spectra, as in n°1). CTX imaging also reveals wrinkle ridge morphologies over green spots, suggesting that lava flows may have covered ancient clay-rich areas between Oxia Planum and Mawrth Vallis. Layered crater walls, observed with HiRISE, seem to confirm the presence of clays between these two major regions – more closely related than we initially believed, but showing signs of a different ancient aqueous history. We are now deciphering the stratigraphic relationship between the two types of clays, and broadening the scope of our study to the clay-rich margins of Chryse Planitia.
Figure 2: Mawrth Vallis (blue) and Oxia Planum (red) clays over pyroxenes (green) and OMEGA dust emission (greyscale) [14].
References: [1] Poulet F. et al. (2005) 438(7068), 623-627. [2] Carter J. et al. (2016) 47th LPSC, Abstract #2064. [3] Vago J. L. et al. (2017) Astrobiology, 17(6-7), 471-510. [4] Quantin-Nataf C. (2021) Astrobiology, 21(3), 345-366. [5] Bishop J. (2008) Science, 321(5890), 830-833. [6] Loizeau D. (2007) JGR: Planets, 112(E8). [7] Bibring J.-P. et al. (2004) ESA Spec. Pub., SP-1240. [8] Jakimow B. et al. (2023) SoftwareX, 23, 101507. [9] Langevin Y. et al. (2005) Science,307(5715), 1584-1586. [10] Jouglet D. et al. (2007) JGR :Planets, 112(E8). [11] Carter J. et al. (2013) PSS, 76, 53-67. [12] Ody A. et al. (2012) JGR: Planets, 117(E11). [13] Carter J. et al. (2023) Icarus, 389, 115164. [14] Audouard J. et al. (2014) Icarus, 233, 194-213.
How to cite: Torres, I., Carter, J., Quantin-Nataf, C., Volat, M., Millot, C., and Dehouck, E.: Extent and Nature of Clay-rich Deposits, from Oxia Planum to Mawrth Vallis, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-609, https://doi.org/10.5194/epsc2024-609, 2024.
Introduction: Hellas Planitia is the largest and deepest preserved impact basin on Mars. It was likely in place by the Early Noachian and has since undergone a complex geologic history with km-thick infillings and similar magnitudes of erosion [1]. As a deep basin, Hellas must have played an important role within the Noachian hydrological system(s), regardless of if these were long or short-lived, continuous or intermittent [e.g. 2,3]. There are a number of conspicuous morphological and mineralogical clues within or around Hellas which suggest past aqueous activity. In particular, its outer rim features phyllosilicates in layered (likely sedimentary) deposits [4,5]. There is also evidence for at least ephemeral flow of water, as evidenced by regionally extensive silcrete deposits [6-9]. Yet there is no evidence for past ponding within Hellas.
Method: To better understand the role of Hellas within past hydrological systems, we conducted a systematic survey of the mineralogy and morphology. Mineralogy is obtained through orbital remote sensing with the OMEGA/Mars Express and CRISM/MRO instruments. Most observations did not yield meaningful surface composition data, owing to the increased atmospheric opacity of Hellas, its frequent hazes and frosts. We focused our study on the most eroded western part of Hellas which exhibits terrains of a large range of ages.
Results: The resulting mineral mapping is shown in figure 1 (MOCAAS database [10] https://www.ias.u-psud.fr/moccas). Three distinct mineral units are distinguishable: 1) Fe/Mg phyllosilicate sediments (red) just outside of Hellas, 2) opal CT/chalcedony silcretes (cyan) on the intermediate margin of Hellas, 3) a previously unknown sulfate-rich floor unit (green). The patchy nature of the mineral maps within Hellas is a detection bias as they are only few valid observations, and signal-to-noise levels are generally lower than the rest of Mars. A conservative mapping approach was chosen here.
We focus on the sulfate unit of the floor of Hellas, which is mostly exposed on its western side. The striking feature of these salts is that they correlate with a small set of morphological units, and that they are found in the deepest exposed sections of Hellas Planitia (currently mapped between - 6500 and -7500 m, mostly below -7000 m). As such they are likely not related to polar sulfates which are known to have formed (or still form) at high southern latitudes [11].
The precise composition of the salt hasn’t been ascertained yet, but it is most likely a poly-hydrated sulfate, which spectra are similar to salt deposits of the Terby crater floor, but is not unique to these sites. The 1.9 µm band shape is generally inconsistent with gypsum but may include bassanite. The spectra do not exclude hydrated chlorinated salts [12,13]. The salt detections are correlated primarily to the so-called “honeycomb” terrains, and to a lesser degree to the “banded” terrains (figure 2). Previous terrain designation also included “reticulate” terrains [12-14].
Discussion: A recurring interpretation of the honeycomb terrains are that they diapirs which topmost section has been eroded away, exposing the truncated cells [e.g.17-19]. The banded terrains are at times equated to the honeycomb terrains or considered to be a different unit. The formation of diapirs is related to the presence of layers of different densities, the lighter material being buried below a dense material such as lava flows. The diapiric material has been proposed to be either salt or ice [e.g. 19] but direct compositional evidence was lacking to distinguish between both scenarios. We therefore propose that these were indeed salt diapirs. The salty nature and cell size of the diapiric terrains of Hellas Planitia suggest ~2 km of salts existed on the basin floor [e.g. [19], only exposed where erosion most deeply incised the overlying material.
One consequence is that this observation may help solve the conundrum of the missing salts on early Mars [20]. As the Noachian highland terrains were leached and formed phyllosilicates, if sulfates were the cation host, it may indicate the climate at least intermittently allowed a mature hydrological system wherein salts would be transported in dissolved fraction over large distances from source to sink, and accumulate in sedimentary basins - most notably Hellas.
Concerning Hellas Planitia specifically, this detection suggests that Hellas has had one or several episodes during which it was filled by water. Kieserite (mono-hydrated sulfate) is relatively dense and would preclude the formation of diapir cells but poly-hydrated salts are light enough to enable diapirs [19], consistent with the CRISM and OMEGA spectra. Chlorinated salts are also light enough and cannot be excluded from the spectra at this point, either as an alternative to the sulfate, or perhaps as an assemblage which cannot be directly detected (especially in dehydrated form).
Further investigations of the extent of the sulfate unit is required to understand both the altitude range and lateral extent of the salt diapirs, which will strongly affect the size of the sea and the cumulative GEL of water required to deposit the salts. Recent estimates range from a sea a few 100s m deep to over 1 km [17], but previous studies proposed Hellas was filled well above these values [4]. The exact sulfate composition will also be refined.
Figure 1 . Western Hellas Planitia aqueous mineralogy based on OMEGA and CRISM. Red: Fe/Mg phyllosilicates. Cyan: hydrated silica. Green: poly hydrated sulfates.
Figure 2. Close up on the sulfate-rich diapirs of Hellas. Geological units from [14,16].
References: [1] Tanaka+95, JGR, 100, 5407-5432. [2] Andrews-Hanna+11, JGR, 116, E02007. [3] Hiatt+24, Icarus, 408, 115774. [4] Wilson+07, JGR, 112, E08009. [5] Ansan+11, Icarus, 211, 273-304. [6] Bibring+06, Science, 312, 400. [7] Bandfield+08, GRL, 35, L12205. [8] Carter+13, JGR, 118, 831–858. [9] Pineau+23, LPSC 54th, abstract #1982. [10] Carter+23, Icarus, 389, 115164. [11] Barraud+22, EPSC 16th, abstract #847. [12] Hanley+13, JGR, 20, 1415–1426. [13] Bishop+14, LPSC 45th, abstract #2145. [14] Leonard+01, USGS Spec. Pub., doi. 10.3133/i2694. [15] Bernhardt+16, Icarus, 264, 407-442. [16] Bernhardt+19, Icarus, 321, 171-188. [17] Mangold+03, 6th Int. Conf. Mars, Abstract #3047. [18] Bernhardt+16, JGR, 121, 714-738. [19] Weiss+17, Icarus, 284, 249 -263. [20] Millliken+09, GRL, 36, L11202.
How to cite: Carter, J., Ansan, V., and Mangold, N.: Was Hellas Planitia a Noachian Sea?, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-802, https://doi.org/10.5194/epsc2024-802, 2024.
Introduction. Visible and near-infrared observations of Mars have revealed a wide range of primary and secondary alteration minerals. Hydrated minerals were first identified from orbit with the Observatoire pour la Minéralogie, l’Eau, les Glaces, et l’Activité (OMEGA) instrument. OMEGA data were used to produce global maps of aqueous minerals [1], including the polar regions [5], in particular the northern one [2]. Recently, the Mars Orbital Catalog of Aqueous Alteration Signatures project [3] including OMEGA observations and data from the Compact Reconnaissance Imaging Spectrometer for Mars (CRISM) produced a set of the more accurate and resolved maps of aqueous minerals at latitudes ranging from 70°S to 70°N. However, no recent systematic study of the distribution and nature of the aqueous composition has been carried out in the south polar regions. OMEGA observations highlighted spectral variability of the Martian high latitude surfaces [5, 4]. In particular, the 1.9 µm and 3 µm band depths increase from 60° latitudes towards the poles [1, 5], which suggests that hydrated minerals may participate in this spectral variability. [4] indeed reported the detection of a deep and relatively narrow absorption band at 3 µm coupled with a shallow but significantly wider 4 µm feature in the northern high latitudes of Mars. This feature may be related to sulfates and has no significant counterpart in the south for now. Moreover, the 1.9 µm band is significantly broader and shifts to slightly longer wavelengths in the south polar region which seems to indicate a specific mineralogy, and/or a different hydration state from that of the north [5]. Here we present a spectral analysis and the detection of hydrated minerals in the south polar region.
Method. OMEGA has been observing the Martian surface since 2004 between 0.35 and 5.1 µm. Here, we selected observations acquired during southern spring and summer as high latitudes are covered by seasonal ice during winter. OMEGA data have been processed using the OMEGA-Py Python module [8]. Our analyses focus on the detection and characterization of hydrated minerals using several spectral criteria previously defined: the band depths (BD) at 1.9 µm [4], 2.1 µm [6] and 2.4 µm [4]; the narrow and wide 3 µm band depths [4] and the SINDEX [6]. We have developed a new estimator adapted to the broader and shifted 1.9 band observed in the southern high latitudes, called here the “wide 2 µm band depth (WBD)”. This criterion estimates the band depth centered at 2 µm using the average continuum from 1.75 to 1.85 µm and 2.20 to 2.30 µm. Since the 3 µm band is very sensitive to the presence of water ice at the surface or in the atmosphere [7], we remove pixels with a water ice absorption at 1.5 µm>1%. In order to prevent possible effects related to viewing geometry, we removed pixels with emergence and incidence angles respectively higher than 10° and 75°. We then isolate detections of sulfate-bearing units thanks to a combination of four of these spectral criteria: the BD at 1.9, 2.1 and 2.4 µm and the SINDEX.
Results. The spatial distribution of all the spectral criteria over the whole OMEGA dataset in the south polar region have been investigated. All the criteria exhibit a strong variability in this region. The narrow 3 μm BD identified for the first time by [4] in the north polar region shows a different behavior in the south polar region. This spectral parameter varies mainly from 0 to 10% in the southern polar region while it increases to more than 30% in the northern polar region. The WBD at 2 µm is present at all longitudes between 50°S and 90°S and varies from 1% to 6% (Fig 1). The wide 3 μm BD is almost always higher than 35% and increases over 65% in this region. The WBD at 2 µm and the WBD at 3 µm increase with latitude toward the south pole (Fig 1). In addition to latitudinal variations, the WBD at 2 µm exhibits seasonal variations and decreases over the summer. The map of hydrated minerals detections in the south polar region (Fig 2) shows a high concentration of sulfate-bearing units around the polar cap. High-resolution CRISM observations obtained over this unit confirm that the spectral signatures can be confidently attributed to sulfates.
Perspectives. The spectral signatures between the sulfate detections presented here and the sulfates in Meridiani Planum are different especially around 3 μm which may indicate different sulfates compositions and/or different mixings with other minerals. Such differences may provide clues about the specific sulfates formation and transformation pathways expected at high latitudes. The ring-like structure of the south polar sulfates suggest a formation associated with seasonal frost and/or polar ice cap. The absence of the narrow 3 μm band depth around the south pole differentiate the two Martian poles in term of spectral variations. Further study of the spectral signatures and comparison with geological maps is in progress, and will help constrain the scenario for the formation of these hydrated minerals in the South Polar region.
References
[1] Poulet et al., (2007). JGR: Planets, 112(E8). [2] Langevin et al., (2005). Science, 307, 5715. [3] Carter et al., (2023). Icarus, 389, 115164. [4] Stcherbinine et al., Icarus, 369, 114627 (2021). [5] Poulet et al., GRL, 35(20) (2008). [6] Viviano-Beck et al., JGR:Planets, 119(6), 1403-1431(2014). [7] Jouglet et al., JGR:Planets 112.E8 (2007). [8] Stcherbinine (2023). Zenodo, doi:10.5281/zenodo.10035061
Figure 1: Composite maps of 833 OMEGA observations in the south polar region of Mars. Solar longitude of the observations ranges from 260° to 340°, corresponding to late southern spring and southern summer. Color codes shows the spectral variations of the wide 2 µm band depth (left) and the wide 3 µm band depth (right).
Figure 2: Map of hydrated mineral detections in the south polar region of Mars (50°S-90°S). Pixels in blue represents monohydrated sulfates detection, in green polyhydrated sulfates detection and pixels in red are phyllosilicates detection. Based map is the wide 3 µm band depth.
How to cite: Barraud, O., Carter, J., Vincendon, M., Stcherbinine, A., and Sheppard, R.: Spectral variability in the south polar region of Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-921, https://doi.org/10.5194/epsc2024-921, 2024.
The Curiosity rover of the Mars Science Laboratory mission is exploring Gale crater on Mars since August 2012 [Vasavada2022]. On board, the ChemCam instrument [Wiens2012, Maurice2012] combines Laser-Induced Breakdown Spectroscopy (LIBS), passive spectroscopy and imaging.
ChemCam is a remote sensing instrument, assessing the chemical composition of rocks and soils up to a few meters from the rover, imaging outcrops at longer distances, and surveying the atmospheric variability. The chemistry of primary and secondary phases is obtained using LIBS by performing typically five point rasters on the martian targets. Each point is investigated with 30 laser shots. This technique can therefore highlight some elemental correlations [Rammelkamp2024], or even some depth variations (<1mm) [Maurice2016]. As of Sol 4183 (martian days), ChemCam has acquired thousands (32,776) LIBS observation points (Fig. 1), revealing the chemostratigraphy of the central mound, Aeolis Mons, in Gale crater, as the rover continues its ascent. More than 41,000 passive spectra and 23,700 images were also collected with ChemCam.
Figure 1: Cumulative number of LIBS points ordered in time, plotted as a function of the rover elevation in the different geologic formations.
One of the main reasons for selecting Gale Crater as the landing site was the orbital observations of hydrated clay minerals on the most basal layers of Aeolis Mons, overlain by hydrated sulfates up-section. This succession of mineral signatures is hypothesized to be a feature of the global-scale climate change that Mars has encountered [Grotzinger2015]. The rover explored the transition between the clay-bearing unit and the sulfate-bearing unit, since Sol 3072 at the Mount Mercou outcrop, and extended up to Sol 3655, in the Marker Band Valley [Edgar2024]. As part of its fourth extended mission, Curiosity’s payload is being used to characterize the salt minerals detected in this unit, seek their origins, and test the hypothesis of a link with the global evolution towards an arid climate [Fraeman2021].
The transition zone can be characterized in different ways. From the point of view of sedimentary facies, it is a transition from lacustrine deposits to eolian deposits [Edgar2024]. From a mineralogical point of view, this zone is characterized by the disappearance of clays (evidence of the presence of liquid water on the surface), in favor of sulfates associated with a more arid environment [Rampe2023]. However, before being found in the bedrock, these sulfates were concentrated in abundant secondary phases (nodules, veins) associated with fluid circulations. ChemCam showed that these sulfates were mainly Mg-rich sulfates [Frydenvang2024, Rapin2024]. On the other hand, Ca sulfates are found mainly in the bedrock even though they have also been observed in nodules. This illustrates the complex interactions between primary deposits, their alteration and later diagenetic phenomena when attempting to reconstruct the history of this transitional period. The observations of centimetric polygonal structures in this area, due to their fragility, has led to the suggestion of a relatively rapid repeated succession between wet and arid phases. Multiple climate transitions [Kite2024] or cyclic climatic variations [Rapin2023] would have implications for exobiology research strategies at the boundary between the Noachian and Hesperian epochs, as the rapid and prolonged succession of wet and dry environments could be conducive to the polymerization of organic molecules in a pre-biotic chemistry. In line with these implications, ChemCam has detected much more apparent halite in this transition zone, with unusual facies compared to previous ones [Meslin2024].
Other intriguing zones have been traversed, such as the Amapari Marker Band, which evokes an unexpected coastline interrupting the otherwise arid depositional environment. All in all, water seems to have played a more important role than expected in this unit, with potential implications for its habitability. For example, ChemCam has detected chemical correlations between Mn and Fe contents, and both are enriched compared to average. In some cases, Zn is also elevated, along with Cu [Gasda2024].
In the overlying sulfate-bearing unit, where the rover is now, alternating light- and dark-toned bands are identified from orbit [Sheppard2021]. From its in situ vantage point, rover observations suggest that distinct dark-toned horizons are found predominantly in the orbitally defined dark-toned bands. Geochemical data from ChemCam on these dark-toned horizons (Fig. 2) reveal they are enriched in fluorite [Forni2024], as well as Na-Mg sulfate phase [Hughes2024]. These observations are unique, and along with the observation of siderite [Tutolo2024], suggest complex brine systems possibly at high temperatures [Forni2024].
Figure 2: Chemistry from a dark-toned horizon (Tenderfoot_Peak_ccam, red points) and light-toned bedrock (Muro_Blanco, blue points) are highlighted, indicating presence of Na-Mg-sulfates in dark material, Mg-sulfates in light material, and broad contributions from Ca-sulfates and Mg-sulfates to the regional geochemistry. Geochemical data from ChemCam in the Sulfate Unit (sols 3949 – 4107). Background MastCam image credits: NASA/JPL-Caltech/MSSS.
The acquisitions of long-distance ChemCam images of the Gediz Vallis ridge show a clast-supported outcrop consisting of blocks with a wide range of textures, colors (provided by MastCam [Bell2017]), and geomorphological signatures. While still under discussion, this could be interpreted as a rock avalanche [Dietrich2024].
References: Bell et al. (2017) Earth and Space Science, 4 (7); Dietrich et al. (2024) 10th Int. Conf. Mars ; Edgar et al. (2024) LPSC, #1016. Forni et al. (2024) 10th Int. Conf. Mars ; Fraeman et al. (2021) Mars Geological Enigmas, Chapter 1 ; Frydenvang et al. (2024) 10th Int. Conf. Mars ; Gasda et al. (2024) 10th Int. Conf. Mars ; Grotzinger et al. (2015) Science, 350 (6257) ; Hughes et al. (2024) LPSC, #2288 ; Kite and Conway (2024) Nature Geoscience 17 (1); Maurice et al. (2012) Space Sci. Rev, 170; Maurice et al. (2016) J. Anal. Atomic Spec., 31 (4); Meslin et al. (2024) 10th Int. Conf. Mars ; Rammelkamp et al. (2024) EPSC, #869; Rampe et al. (2023) LPSC, #1554; Rapin et al. (2023) Nature, 620 ; Rapin et al. (2024) 10th Int. Conf. Mars ; Sheppard et al. (2021) J. Geophys. Res., 126 (2) ; Tutolo et al. (2024) LPSC, #1564 ; Vasavada et al. (2022) Space Sci. Rev. 218; Wiens et al. (2012) Space Sci. Rev., 170.
How to cite: Gasnault, O., Schroeder, S., Cousin, A., Frydenvang, J., Hughes, E., Dehouck, E., Rammelkamp, K., Johnson, J., Rapin, W., Lanza, N., Wiens, R. C., and Maurice, S. and the ChemCam science team: Exploring the sulfate-bearing unit: Recent ChemCam results at Gale crater, Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1095, https://doi.org/10.5194/epsc2024-1095, 2024.
Clouds play a crucial role in the past and current climate of Mars. Cloud particles impact the planet's energy balance and atmospheric dynamics, as well as influence the vertical distribution of dust particles through dust scavenging. This process of dust scavenging by clouds has significant consequences for the planet's water cycle. For example, regions in the atmosphere with insufficient quantities of dust particles, or condensation nuclei, can inhibit the formation of H2O clouds, leading to the presence of water vapor in excess of saturation [1]. Recent observations made by the MEDA Radiation and Dust Sensor (RDS) [2,3] have shown a marked decline in mesospheric cloud activity (above 35-40 km) when Mars is near its aphelion (within the Aphelion Cloud Belt-ACB season), notably occurring during solar longitudes (Ls) between Ls 70° and 80° [4] (see Figure 1).
In order to investigate the possible factors leading to this decrease in water ice abundance, we used a one-dimensional cloud microphysical model [5,6], which includes the processes of nucleation, condensation, coagulation, evaporation, precipitation, and coalescence, and where the vertical mixing is parameterized using an eddy diffusion profile (Keddy). Combining cloud microphysics modeling with ground-based (Mars 2020 and InSight) and orbital observations (TGO and MRO) of clouds, water vapor, and temperature, we will discuss in this presentation the main factors controlling the water abundance in the Martian mesosphere during the ACB season.
References: [1] Maltagliati, Luca, et al. "Evidence of water vapor in excess of saturation in the atmosphere of Mars." science 333.6051 (2011): 1868-1871. [2] Apestigue, V., et al. “Radiation and Dust Sensor for Mars Environmental Dynamic Analyzer Onboard M2020 Rover”. Sensor 22.8 (2022): 2907. [3] Rodriguez-Manfredi, Jose Antonio, et al. “The Mars Enviromental Dynamics Analyzer, MEDA. Asuite of enviromental sensors for the Mars 2020 mission.” Space science reviews 217.3 (2021): 1-86. [4] Toledo, D., et al. “Measurement of aerosol optical depth and sub-visual cloud detection using the optical depth sensor (ODS)”. Atmospheric Measurement Techniques 9.2 (2016): 455-467. [5] Montmessin, F., Rannou, P., Cabane, M.: New insights into martian dust distribution and water-ice cloud microphysics. Journal of Geophysical Research: Planets 107(E6), 41 (2002). [6] Rannou, P., Montmessin, F., Hourdin, F., Lebonnois, S.: The latitudinal distribution of clouds on titan. science 311(5758), 201205 (2006).
How to cite: Toledo, D., Rannou, P., Apestigue, V., Rodriguez-Veloso, R., Arruego, I., Martinez, G., Tamppari, L., Munguira, A., Lorenz, R., Stcherbinine, A., Montmessin, F., Sanchez-Lavega, A., Patel, P., Smith, M., Lemmon, M., Vicente-Retortillo, A., Newman, C., Viudez-Moreiras, D., Hueso, R., and Bertrand, T. and the +4: Decline in Water Ice Abundance in the Martian Mesosphere during Aphelion, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-504, https://doi.org/10.5194/epsc2024-504, 2024.
Thanks to successful long-lasting missions (e.g. Mars Express) and the launch of new space missions (e.g. ExoMars), the amount of Mars’ atmosphere remote sensing observations currently available has significantly increased, allowing to study the planet in greater details than ever before. The methods used to analyze these space observations are generally based on iterative approaches, for which the radiative transfer is solved at each iteration and can become time-consuming when applied to large datasets. This has motivated us to try applying an alternative method developed for the detection and quantification of trace gases in the Earth’s atmosphere, and which will be presented. This detection method is based on the calculation of a spectrally integrated index, called “hyperspectral range index" (HRI), which represents the strength of the spectral signature of the targeted atmospheric species in a spectral range. The HRI yields a dimensionless scalar and is calculated as
with K the Jacobian of the target species, Sy a covariance matrix generated from a representative set of spectra that do not include the target species spectral signature, and ȳ the mean spectrum of the covariance matrix. In that context, Sy is a statistical characterization of the expected correlations between the spectral channels in the absence of the target species. The HRI can encompass broad spectral ranges to exploit all the channels in which the target species is absorbing, which results in a substantial gain of sensitivity over other detection methods and makes it suitable for the detection of broadband absorption features. As the HRI can lead to false detections in case of partial match between spectral signatures of the target species and an interference, a firm identification of the target species is needed to confirm its contribution to the HRI. To do so, a transformation, referred to as whitening, has been applied.
The method has been tested on nadir observations of the PFS (Planetary Fourier Spectrometer, Giuranna et al., 2021) instrument onboard Mars Express. The spectral range 350-500 cm-1has been exploited to calculate the H2O vapor HRI. The covariance matrix and its associated mean spectrum have been built using spectra recorded in regions known to be depleted in water vapor. Using these, the HRI has then been calculated for the entire PFS dataset. Seasonal distributions have been built (an example is shown on Figure 1 for the Martian Year 27) and will be discussed.
As the HRI value depends on the H2O abundance in Mars’ atmosphere, it is possible to convert it to an integrated column. For this, we have built an artificial feedforward neuronal network (NN) that will be presented. The use of a NN allows us to take into account the different parameters affecting the HRI value in addition to the water column, such as the temperature profile and the surface temperature. These parameters are included as input parameters of the NN. The NN has been trained with more than 200,000 forward-simulated PFS spectra with the line-by-line Atmosphit radiative-transfer code (Coheur et al., 2005) using PFS-L2 auxiliary data, which values are randomly selected to be fully representative of real Martian conditions. Still under analysis, the NN yields a target gas total column per single PFS measurement, for which the uncertainties are estimated by propagating the uncertainties associated with each input variable through the NN.
How to cite: Es-Sayeh, M., Bauduin, S., Clarisse, L., Franco, B., and Giuranna, M.: First Planetary Science Application of the Hyperspectral Range Index: Water Vapor as Observed by PFS/MarsExpress, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1099, https://doi.org/10.5194/epsc2024-1099, 2024.
Introduction: We have currently accumulated over 25 years of continuous satellite data on Martian dust and, generally, on the weather of the Red Planet. By utilizing data from instruments operating in the thermal infrared such as the Thermal Emission Spectrometer (TES, onboard the Mars Global Surveyor satellite), Thermal Emission Imaging System (THEMIS, onboard Mars Odyssey), and Mars Climate Sounder (MCS, onboard the Mars Reconnaissance Orbiter), we have been able to reconstruct diurnal maps of column dust optical depth (CDOD) spanning more than 13 Martian years (MY) from 1999 to the present day [1, 2, 3, 4, 5]. Two types of maps exist: ‘gridded maps’ (the mesh is regular but values are missing where there are no observations) and ‘kriged maps’ (interpolated from the gridded maps using kriging to produce complete maps). These longitude-latitude maps are used as 'dust scenario' in the Mars Climate Database [6], among many other applications. They are routinely updated and made publicly available in NetCDF or FITS formats (see the links to the datasets in the acknowledgments section).
Recent developments: The daily CDOD maps covering MY 24 through part of MY 27 have been recently improved by using revised retrievals of column dust optical depths from TES observations (refer to the link in the acknowledgments section for accessing this updated dataset). The top four panels of Figure 1 (MY 24, 25, 26, 27) show zonal means of these improved CDOD maps normalized to 610 Pa (MY 27 zonal mean is actually a combination of revised TES-based maps and previous THEMIS-based maps).
The arrival of the Emirates Mars Mission (EMM) 'Hope' spacecraft in a low-inclination, high-altitude orbit around Mars has enabled simultaneous monitoring of the full disk of the Martian atmosphere. The CDOD retrievals from the Emirates Mars InfraRed Spectrometer (EMIRS) observations significantly enhance the quality of our dust maps, enabling for the first time quasi-continuous monitoring of storms over multiple local times [7]. The bottom two panels of Figure 1 show zonal means of, respectively, MCS-based maps and EMIRS-based maps for MY 36. Work is in progress to 1) understand the differences, and 2) integrate CDOD information from both MCS and EMIRS to produce combined daily maps.
Moreover, the availability of visible images from the EMM/Emirates Exploration Imager (EXI) on one side [8], and retrievals of CDOD in the visible from TES Emission Phase Function (EPF) observations on the other (refer to the link in the acknowledgments section for access to this novel dataset), enables cross-comparison and validation of the daily gridded dust maps with an unprecedented level of detail.
Figure 1: Zonal means of CDOD normalized to 610 Pa for 13 Martian years. The top four panels (MY 24, 25, 26, 27) show zonal means of improved maps using revised TES retrievals. The bottom two panels show zonal means of, respectively, MCS-based maps and EMIRS-based maps for MY 36.
Tracking of dust events: From the daily CDOD maps, it is possible to identify large-scale dust events (“storms”) reaching regional and planetary scales, follow their evolution, and create statistics of their main characteristics such as trajectory, area, and optical depth (see an example in Figures 2 to 5). A new development in progress is the identification and tracking of large-scale dust events using unsupervised machine learning algorithms [9].
A key outcome of this work is the production of a catalog of historic large-scale dust events, which can be routinely updated with new events as new dust maps become available. An important aspect of creating a catalog of dust events is the precise determination of their occurring time. While a Mars calendar based on solar longitude works for many applications, it does not work well for daily maps produced at a discrete number of sols per year. Therefore, we choose to use a sol-based calendar as described in [1, Appendix A]. Each Martian year has either 669 or 668 sols, following a 5-year cycle. The beginning and end of a year is always at midnight at the prime meridian, hence a new year in our sol-based calendar does not necessarily start at LS=0° (see Figure 6).
Figure 2: Sol-by-sol identification and tracking of the evolution of a dust sequence (“storm”) in MY 36 between LS=309° and LS=318° from daily EMIRS-based CDOD maps normalized to 610 Pa. The rest of the daily maps is visible in the transparent background.
Figure 3: Plot of the trajectory of the centroid of the dust event shown in Figure 2.
Figure 4: Plot of the time evolution of the area of the dust event shown in Figure 2.
Figure 5: Plot of the time evolution of the average and 1-σ envelope of the CDOD normalized to 610 Pa for the dust event shown in Figure 2.
Figure 6: This table shows the Earth UTC dates and Mars solar longitudes of the beginning of Martian years 1 through 38 in our sol-based calendar (see also [1]).
Acknowledgments: LM acknowledges support from CNES and ESA MCD project. BKG was supported by UAE University Grant G00003407. Work at the Jet Propulsion Laboratory, California Institute of Technology, is supported by NASA.
The multi-annual dataset of daily gridded and kriged maps v2.x for MY24 through MY36 is available on the MCD webpage (NetCDF format) at https://www-mars.lmd.jussieu.fr/mars/dust_climatology/ and on the VESPA repository (FITS format) at https://bit.ly/3QMFfIf (shortened link)
The latest v3.0 of the daily gridded maps for MY24 through part of MY27, together with corresponding TES CDOD retrievals in the infrared and in the visible, are available on the NASA PDS (atmosphere node) at:
https://atmos.nmsu.edu/data_and_services/atmospheres_data/MARS/montabone.html
References:
[1] L. Montabone et al., 2015, doi: 10.1016/j.icarus.2014.12.034.
[2] L. Montabone et al., 2020, doi: 10.1029/2019JE006111.
[3] M. D. Smith, 2009, doi: [10.1016/j.icarus.2009.03.027
[4] M. D. Smith, 2004, doi: 10.1016/j.icarus.2003.09.010
[5] A. Kleinboehl et al., 2009, doi: 10.1029/2009JE003358
[6] E. Millour et al., EPSC2022-786, https://doi.org/10.5194/epsc2022-786, 2022.
[7] Smith et al, 2022, doi : 10.1029/2022GL099636
[8] B. K. Guha et al., 2023, doi: 10.1029/2023JE008156
[9] T. Lombard & L. Montabone, 10th International Conference on Mars, 2024
How to cite: Montabone, L., Edwards, C. S., Forget, F., Kass, D., Kleinboehl, A., Guha, B. K., Guyon, V., Lombard, T., Millour, E., Smith, M. D., and Wolff, M. J.: Monitoring Martian Atmospheric Dust using Multi-annual, Multi-instrument Orbital Observations, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-381, https://doi.org/10.5194/epsc2024-381, 2024.
Aerosols on Mars are a primary elements for studying the interaction between the solar radiation and the atmosphere and surface. Depending on properties such as aerosol number density, particle radius, or refractive index, the impact of the aerosols can provide positive or negative radiative feedbacks on the dynamics of the atmosphere. Previous studies have revealed large temporal and spatial variability in the aerosol optical properties, emphasizing the necessity for continuous monitoring of these properties throughout the day and at multiple locations. To address these measurements, the Radiation and Dust Sensor (RDS) [1] is part of the Mars Environmental Dynamics Analyzer (MEDA) [2] payload onboard of the Mars 2020 rover Perseverance. RDS instrument compromises two sets of 8 photodiodes (RDS-DP) and a camera (RDS-SkyCam). One set of photodiodes is pointed upward, with each one covering a different wavelength range between 190-1200 nm. The other set is pointed sideways, 20 degrees above the horizon, and they are spaced 45 degrees apart in azimuth to sample all directions at a single wavelength. The analysis of these observations with a radiative transfer model [3] (Fig. 1) allow us to fit aerosol parameters such as the aerosol opacity at different wavelengths or the aerosol particle radius. In this work we will discuss some preliminary results for the first 100 sols of Mars 2020 mission.
References:
[1] Apestigue, V., et al. “Radiation and Dust Sensor for Mars Environmental Dynamic Analyzer Onboard M2020 Rover”. Sensor 22.8 (2022): 2907.
[2] Rodriguez-Manfredi, Jose Antonio, et al. “The Mars Enviromental Dynamics Analyzer, MEDA. Asuite of enviromental sensors for the Mars 2020 mission.” Space science reviews 217.3 (2021): 1-86.
[3] Toledo, D., et al. “Measurement of aerosol optical depth and sub-visual cloud detection using the optical depth sensor (ODS)”. Atmospheric Measurement Techniques 9.2 (2016): 455-467.
How to cite: Rodriguez-Veloso, R., Toledo, D., Apéstigue, V., Arruego, I., Lemmon, M. T., Smith, M. D., Martínez, G., Vicente-Retortillo, Á., Jiménez-Martín, J. J., García-Menéndez, E., Viudez-Moreiras, D., Sánchez-Lavega, A., Pérez-Hoyos, S., Sebastian, E., de la Torre-Juárez, M., and Rodríguez-Manfredi, J. A.: Aerosol optical properties observed by MEDA Radiation an Dust Sensor (RDS) at Jezero Crater, Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-508, https://doi.org/10.5194/epsc2024-508, 2024.
The presence of the Martian oxygen green line emission at 557.7 nm in the Martian dayglow was discovered with the NOMAD-UVIS ultraviolet-visible spectrometer on board ESA’s Trace Gas Orbiter (Gérard et al., 2020). It corresponds to the 1S -1D forbidden transition in the O atom. Limb profiles of the emission have been observed with a vertical resolution of about 10 km. They generally show two peaks near 80 km and 110 km during quiet solar periods. Intensity and altitude variations of the emission peaks have been previously reported (Soret et al., 2022). The green line is by far the strongest feature in the Mars dayglow spectrum.
Fig. 1: NOMAD/UVIS dayside limb spectra binned into altitudes regions between 70 and 160 km. In addition to the strong green line, the OI line at 297.2 nm and the CO2+ UV doublet at 288-289 nm are also observed. The spectral resolution varies from 1.2 nm at 200 nm to 1.6 nm at 600 nm (Gérard et al., 2020)
The O(1S) level is mainly excited by photodissociation of CO2 by solar EUV radiation, with minor contributions mostly from electron impact on CO2 molecules and dissociative recombination of CO2+ ions. The doublet at 630-636.4 nm has also been observed but is about 20 times weaker (Gérard et al., 2021; Soret et al., 2022). The lower peak is excited by solar Lyman-a radiation while the upper one is produced by a range of EUV radiation.
Dayside observations have been performed since April 2019 and monitored to observe the changing brightness and peak altitude as solar activity increases. Observations of the UV counterpart 1S-3P at 297.2 nm of the green line have shown that the upper emission peak remains located at a pressure level close to 0.39 microbar (Gkouvelis et al., 2018). The solar EUV flux at the Mars location is provided in three spectral channels by the EUV monitor (EUVM) on board MAVEN. UVIS dayglow observations collected at different seasons and latitudes have shown that the 557.7-nm emission varies with season and latitude (Soret et al., 2022), making it possible to follow CO2 density variations in the mesosphere.
Fig. 2: Limb profiles of the OI 557.7-nm dayglow (c) in northern Spring and Summer, (d) in southern Spring and Summer Numerical simulations for the same conditions are shown in red (Soret et al., 2022).
In this study, we present time variations of the altitude and brightness of the green line distribution during the period of rising activity. We show that the limb brightness directly responds to the intensity of the solar Lyman-alpha flux, while the attitude shows a more complex response dependence, a signature of the expansion and contraction of the atmosphere.
REFERENCES
Gérard, J. C., Aoki, S., Willame, Y., Gkouvelis, L., Depiesse, C., Thomas, I. R., et al. (2020). Detection of green line emission in the dayside atmosphere of Mars from NOMAD-TGO observations. Nature Astronomy, 4(11), 1049–1052.
Gérard, J. C., Aoki, S., Gkouvelis, L., Soret, L., Willame, Y., Thomas, I. R., et al. (2021). First observation of the oxygen 630 nm emission in the Martian dayglow. Geophysical Research Letters, 48(8), e2020GL092334. https://doi.org/10.1029/2020GL092334.
Gkouvelis, L., Gérard, J. C., Ritter, B., Hubert, B., Schneider, N. M., & Jain, S. K. (2018). The O(1S) 297.2-nm dayglow emission: A tracer of CO2 density variations in the martian lower thermosphere. JGR Planets, 123(12), 3119–3132.
Soret, L., Gérard, J.-C., Aoki, S., Gkouvelis, L., Thomas, I. R., Ristic, B., et al. (2022). The Mars oxygen visible dayglow: A Martian year of NOMAD/UVIS observations. JGR Planets, 127, e2022JE007220.
How to cite: Gérard, J.-C., Soret, L., Robin, H., Thomas, I., Ristic, B., Vandaele, A. C., and Hubert, B.: Variations of the Martian oxygen green line dayglow: response to solar activity, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-896, https://doi.org/10.5194/epsc2024-896, 2024.
A compact Raman spectrometer [1] was developed for in-situ science on Phobos with the IDEFIX rover [2], which will be launched as part of JAXA’s Martian Moon eXploration (MMX) mission in 2026. The MMX mission is dedicated to study the two Martian moons, Phobos and Deimos, with the aim of better understanding their origin and evolution [3, 4]. Orbital observations with the main MMX spacecraft and in-situ science on Phobos with the IDEFIX rover will be complemented by returning samples (>10 g) from the surface of Phobos back to Earth. The rover is expected to land between late 2028 and early 2029, in conjunction with the rehearsal for the first landing of the main spacecraft.
Figure 1. left: RAX flight model (FM) during thermal vacuum testing at DLR. Right: RAX development model (DM) combined locomotion test with MMX demonstration rover in DLR testbed during a Raman measurement.
The RAman spectrometer for MMX (RAX) contributes directly to the high level mission objectives by characterizing the mineral composition of the Phobos surface in-situ, providing ground truth. While different mineral phases can be associated with different formation scenarios (e.g. captured asteroid or major impact) and surface alteration processes on Phobos, the mineralogical information obtained with RAX can also help inform decisions about sampling with the main spacecraft and be compared with data obtained from the returned samples and in-situ data from the surface of Mars. The rover's mobility allows measurements to be taken from multiple positions along the rover's traverse, allowing local variations in composition to be studied. RAX measures down to the ground below the rover. RAX is extremely compact, with a mass of only 1.5 kg and a volume of about 1 dm³. The optical design of RAX was driven by 1) the tight volume and mass constraints available of the small MMX rover and 2) optimizing the collection and detection capabilities of the Raman signal from a sample at several centimeters distance below the rover’s body. Raman excitation (λ = 532 nm, typical optical power on sample 20 mW) is provided by a separate laser module based on the Raman Laser Spectrometer (RLS) laser developed for the Rosalind Franklin rover of the ExoMars mission [5]. The laser emission is transmitted to the RAX Spectrometer Module via a multimode optical fiber with a core diameter of 50 µm. To focus the laser on the Phobos surface below the rover, the spectrometer is equipped with an opto-mechanical autofocus subsystem. It allows fine-tuning of the focus position at a working distance of about 8 cm, within a range of 13 mm and an accuracy of 50 µm. The RAX instrument covers a spectral range of 535 to 680 nm, corresponding to a Raman shift of approximately 90 to 4000 cm−1 and therefore enabling the identification of water-bearing minerals. The spectral resolution across the entire spectral range is about 10 cm−1. To demonstrate the functionality of the RAX instrument after launch and to monitor its performance, a Verification Target (VT) is part of the payload. The VT is a 13 mm diameter pellet made of deuterated polyethylene terephthalate (PET) and was developed and space-qualified specifically for this mission [6]. The VT is placed in the field of view of the RAX instrument and can be measured during the cruise until the rover separates from the MMX spacecraft for landing on Phobos.
For a Raman measurement on Phobos, the rover’s locomotion system is needed for the main height adjustment: the rover will gradually lower its body height until the RAX instrument is within its working range. After autofocusing, RAX will measure what is exposed beneath the rover from a footprint of 50 µm. Visual context is provided by the rover’s cameras. To improved data acquisition, Raman data will be collected during Phobos nights.
Representing a collaborative effort, RAX is a joint contribution from the German Aerospace Center (DLR), Instituto Nacional de Técnica Aerospacial (INTA), and JAXA. Delivered to the MMX Phobos Rover in August 2022, the RAX flight model was successfully integrated into the IDEFIX rover in October 2022. The flight model of the rover has been delivered to JAXA/MELCO where it is currently being integrated into the main spacecraft for further qualification and functional testing. Operational sequences, e.g. defining the interaction between locomotion and science instruments, are currently being prepared. RAX scientific measurements are being prepared with laboratory studies using the RAX development model.
Acknowledgements
MMX is a JAXA mission with contributions from NASA, CNES and DLR. The MMX IDEFIX rover is provided by CNES and DLR. The RAX instrument is a joint development from DLR-OS, INTA/UVa, and JAXA/UTo. The authors thank their respective national funding frameworks including DLR’s Programmatik Raumfahrt and PID2022-142490OB-C31 funded by MCIN/AEI/10.13039/501100011033 and are grateful to the MMX, IDEFIX rover and RAX instrument teams who made this work possible.
References
[1] T. Hagelschuer, U. Böttger, M. Buder et al., Internat. Astronautical Congress (IAC): 18-22 September, 2022, Proceed. 2022, IAC-22-A3.4A.8.
[2] S. Ulamec, P. Michel, M. Grott et al., Acta Astronaut. 2023, 210, 95.
[3] T. Usui, K. Bajo, W. Fujiya et al., Space Sci. Rev. 2020, 216, 49.
[4] Y. Kawakatsu, K. Kuramoto, T. Usui et al., Acta Astronaut. 2023, 202, 715.
[5] F. Rull, S. Maurice, I. Hutchinson et al., Astrobiology 2017, 17, 627.
[6] A. G. Moral, J. Mora, O. Prieto-Ballesteros et al., J. Raman Spectrosc. 2023, 54, 1268.
How to cite: Schröder, S., Böttger, U., Cho, Y., Hübers, H.-W., Prieto-Ballesteros, O., Rull, F., Buder, M., Bunduki, Y., Dietz, E., Hagelschuer, T., Kopp, E., Moral Inza, A., Pertenais, M., Rammelkamp, K., Ryan, C., Säuberlich, T., Schrandt, F., Ulamec, S., Usui, T., and Weber, I. and the RAX Team: The Raman Spectrometer RAX on the MMX IDEFIX Rover for in-situ Mineralogical Analysis on Phobos, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-988, https://doi.org/10.5194/epsc2024-988, 2024.
The Curiosity rover, part of the Mars Science Laboratory (MSL), landed in Gale crater in 2012 with the Rover Environmental Monitoring Station (REMS)[1] onboard. REMS includes an atmospheric relative humidity sensor, REMS-H, provided by the Finnish Meteorological Institute (FMI)[2]. REMS-H has continuously recorded hourly near-surface humidity conditions on Mars since landing, resulting in a rich dataset. As of May 2024, REMS-H has been operational for a little over 4100 sols; more than 11 Earth years, providing the longest relative humidity record from the surface of Mars.
Figure 1: REMS-H is located in the boom of Curiosity’s mast. Credit: NASA/JPL-Caltech/MSSS
The current calibration of REMS-H has been evaluated using new calibration measurements performed under a Martian analogue environment at DLR PASLAB (Planetary Analog Simulation Laboratory) at the German Aerospace Center (DLR)[3]. The capacitive sensor type, that has been used in all in situ measurements so far, can be sensitive not only to relative humidity but also other variables like temperature, pressure and carbon dioxide (CO2)[4].
Based on the findings a revised calibration has been developed for REMS-H. This presentation outlines both the revised calibration and the corresponding updated results. Particularly two aspects of the calibration can now be corrected using real measurement data: the RH response function between 0% and 100% rh, and the dynamic range of the sensor in CO2. Furthermore, a two-part calibration correction based on flight data analysis is described.
The revisited results section presents the revised interannual, seasonal, and diurnal variations in relative humidity, temperature, and derived water vapor mixing ratio (VMR). Comparisons with previous calibration results are also discussed. In general, the new calibration resulted in somewhat lower relative humidity values, although the difference varies. On average the difference is about 10% rh. In VMR, the difference is more prominent and temperature dependent. A couple of example sols with current and revised calibration are shown in Figure 2. The full dataset will be available in a subsequent manuscript that is currently under preparation [5].
The resulting new dataset is well aligned with orbital observations and M2020 MEDA HS observations from the same time period. Also the UH/FMI single-column model (SCM)[e.g. 6] has been used to evaluate the revised results with good agreement.
In conclusion, the recalibration effort has improved the accuracy and reliability of REMS-H data, aligning the results with orbital observations and simulation runs. While we believe that the new measurements provide a more accurate representation of humidity values on Mars, it's important to consider the relatively large uncertainty, particularly when RH levels are low, when using the data.
Figure 2: Example comparisons of old (red) and revised (blue) relative humidity measurements from sols 2530 (Ls=81.6°) and 3211 (Ls=87.2°) during southern winter. Sensor temperature is plotted in orange. In most sols the revised RH is lower than previously, but that is not always the case as can be seen in sol 3211 night time observations.
References:
[1] J. Gomez-Elvira et al.: REMS: The Environmental Sensor Suite for the Mars Science Laboratory Rover. Space Science Reviews 170, 583–640, 2012.
[2] Harri, A.-M. et al.: Mars Science Laboratory relative humidity observations: Initial results. Journal of Geophysical Research: Planets, 119, 2132–2147, 2014.
[3] Hieta et al.: Improving relative humidity measurements on Mars: New laboratory calibration measurements. Geoscientific Instrumentation, Methods and Data Systems, in review
[4] Lorek, A. and Majewski, J.: Humidity Measurement in Carbon Dioxide with Capacitive Humidity Sensors at Low Temperature and Pressure. Sensors, 18, 2615, 2018.
[5] Hieta et al.: REMS-H Revisited: Updated calibration and results of the humidity sensor of the MSL Curiosity. Space Science Reviews, In preparation
[6] Savijärvi, H. I., Harri, A.-M., & Kemppinen, O.: The diurnal water cycle at Curiosity: Role of exchange with the regolith. Icarus, 265, 63–69, 2016.
How to cite: Hieta, M., Jaakonaho, I., Polkko, J., Genzer, M., Savijärvi, H., Harri, A.-M., Lorek, A., Garland, S., de Vera, J.-P., Martínez, G., Fischer, E., Sebastián Martínez, E., Rodríguez-Manfredi, J. A., Tamppari, L., de la Torre Juárez, M., and McConnochie, T.: Updated calibration and results of the REMS-H humidity sensor of the MSL Curiosity, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-774, https://doi.org/10.5194/epsc2024-774, 2024.
Introduction
The two martian moons, Phobos and Deimos, will be visited by the sample return JAXA mission MMX in 2027 (launch planned in October 2026). They orbits Mars at respectively 9400 and Deimos at 23000km, with low inclinations and eccentricities. Their heavily processed surfaces harbour a fine regolith [1]. Visible and Near-Infrared (Vis-NIR) spectral readings reveal dark and red flat spectra[2,3]. These characteristics have spawned competing hypotheses regarding their formation:
- They could be remnants of D-type asteroids, captured by Mars[3].
- They may have formed from the aftermath of an impact between Mars and a protoplanet[4].
While the former hypothesis aligns with spectral similarities observed in D-type asteroids, it fails in explaining the moons' peculiar orbits. Conversely, the latter resolves the orbital puzzle but fails in elucidating their spectral characteristics.
To unravel the origins of Phobos and Deimos, the Martian Moons Explorer (MMX) mission[5] aims to scrutinise the moons' compositions. The MMX Infrared Spectrometer (MIRS) will observe the two moons in the range of 0.9 to 3.6µm[6]. However, at a distance of 1.5AU, the moons' surfaces emit a significant thermal flux, especially beyond 2µm, necessitating correction in spectral measurements.
We present a model adapted from previous works[7] to characterise the thermal behaviour of the Martian moons in preparation of MIRS data interpretation. From their properties—such as albedo, inertia, and thermal conductivity—alongside the incoming flux, we aim to compute their surface temperatures accurately, aiding in the interpretation of MIRS spectral observations.
Illumination at the moons’ surface
Given the absence of an atmosphere and the moons small size, their surface temperatures are dictated by the absorbed flux. Consequently, the initial step in computing their physical surface temperatures involves discerning the incident flux.
We developed a model that computes the flux reaching the surface of the moons over one orbit around the Sun. The model relies on the SPICE/NAIF toolkit to computes the distance and the viewing incidence and reflection angles at the moons surfaces. It takes into account for the eclipses when Phobos or Deimos enters Mars' shadow cone, and for the reflected and emitted light on Mars in the Vis-NIR.
Additionally, our model incorporates the shape model provided by[8] to compute illumination angles on the moons' surfaces. Furthermore, it utilizes the Mars albedo map from[9].
In the model, the incidence flux at a specific point of the surface hence varies over a year with:
-solar distance,
-eclipses,
-moon orientation,
-Mars thermal emission.
Fig.1: Incident flux at Phobos surface on the 28/07/2005.
Fig.1 illustrates the incident flux map at Phobos surface at a specific date (here on the 28/07/2005 at 09:27).
It clearly displays the sub-solar point at 175°E, and the sub-Mars point (at 0°E) where a second peak of flux arises as a consequence of the reflected and emitted light on Mars (resp. ~26W.m-2$ and ~11W.m-2).
Solving heat equations
To retrieve the surface temperature, we adapted a thermal model [7] to Phobos and Deimos.
This model solves the 1-D time-dependent heat equation for each point of the incident flux map previously computed.
It takes as input the thermal properties of the surface:
-thermal inertia
-porosity
-bulk density
It does not take into account for convection, but considering the absence of atmosphere the model is still applicable. However, it does not consider radiation, which may reveal itself important for high porosity at the surface, where conduction is reduced to punctual contacts between grains, giving more importance to heat transfer by radiation. Fig.2 presents a temperature overview at Phobos surface, computed from the incident flux map shown in Fig.1, for an inertia of 100J.m-2.s-1/2.K-1 a porosity of 0.9 and a density of 1300kg.m-3. The temperature range computed for the period of observation is similar to the results of [10]. The sub-solar point is clearly visible and reaches 295K, whereas the boreal pole is much cooler at 140K.
In this simulation, the influence of Mars was taken into account.
Fig.2: Phobos temperature map surface the same day.
Influence of Mars: the case of Phobos
Computing the surface temperature with and without Mars, allow to determine its influence on Phobos temperature.
Fig.3 illustrates the ΔT-map on Phobos surface on the hemisphere pointing directly to Mars. It results in an increase of almost 3.5K of the surface temperature at the sub-Mars point. This effect is far from negligible, as even this slight temperature variation lead to a significant increase of 26% in the flux emitted by Phobos' surface at 2.5µm, and may be even more important in some configurations where the polar region will be illuminated only by Mars.
Fig.3: Mars contribution on Phobos surface temperature, up to 3.5K.
Perspectives
We show the importance of a model that takes into account for Mars’s reflection and emission.
However, this model is still under development and currently does not consider the infrared self-heating of the moons. This factor could significantly alter temperatures within craters, which might be crucial for correcting MIRS spectra, particularly in heavily cratered regions.
The temporal albedo variations of Mars (i.e. the dust storms) will also be added in the model as they will completely remove the thermal emission of Mars surface, and increase the albedo.
Acknowledgments
This work is funded by the DIM ORIGIN (Ile de France).
References
[1] Thomas, 1979
[2] Fraeman+, 2012
[3] Fraeman+, 2014
[4] Craddock, 2011
[5] Kuramoto+, 2022
[6] Barucci+, 2021
[7] Ferrari & Leyrat, 2006
[8] Willner+, 2014
[9] Christensen+, 2001
[10] Kuzmin & Zabalueva, 2003
How to cite: Sultana, R., Barucci, A., Leyrat, C., and David, G.: Modelling Martian Moons' Surface Temperature, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-237, https://doi.org/10.5194/epsc2024-237, 2024.
Introduction
Salt flats are deposits originated from the evaporation of mineral-rich bodies of water, these deposits are usually located in places with growing aridity, which results in the accumulation of chlorides where water has disappeared [1]. These deposits are usually found in the bottom of dried or ephemeral lakes, where the evaporation of water exceeds the recharge, and chlorides are precipitated. Salt flats are important planetary analogues due to their extreme environmental conditions and their resemblance to the processes occurred in other planetary surfaces, especially on Mars [2].
[3] identified several outcrops of chlorides in the martian surface using data from the Thermal Emission Imaging System (THEMIS), on board Mars Odyssey [4]. Most of these locations have been covered by high resolutions datasets, but detailed analysis of most of them remains to be done. In this work we compare one of these locations, near a crater system in west Icaria Planum (-44.22°, -125.99°), with the Uyuni salt flat in Bolivia (-20°, -67°), the largest deposits of this type on Earth, to evaluate the potential of Uyuni as an environmental analogue of Mars.
Data and Methods
For Mars, we used four types of datasets. For the visible range we utilized the global CTX mosaic and two panchromatic and colour HiRISE images [5]. We used two hyperspectral CRISM cubes to check the mineralogy [6]. We used the global MOLA DEM for terrain information (Smith et al., 2001). Finally, we checked the global THEMIS day and night global mosaic [4]. For Earth, we used Landsat-8 images [7].
To analyse the geological setting and evolution of Icaria Planum we used a hybrid mapping approach, incorporating geomorphological and spectral information in a single product [8]. We first created geomorphological units, which were later refined or modified according to the spectral information. We then compared our results with satellite images and geological maps of the Uyuni salt flat. The analysis and mapping of the data were done in the geoprocessing software QGIS with the aid of the Mappy plug-in [9].
Results and discussion
The chloride deposits in Icaria Planum are located on top of a wide and flat area, the zone is a volcanic plateau of late Noachian age according to the global map of [10]. The plateau is surrounded by impact craters, but it is the lowest terrain in the surroundings. Spectral signatures from the CRISM cubes confirm the presence of chlorides in the northern part of the plateau (Figure 1). They mainly appear as unaggregated blocks scattered in the plateau, only some spots show a continuous slab of material, which implies the material has been heavily reworked (Figure 2). Between the chloride deposits are some field of mafic dunes, which are the result of the erosion of the original lavas materials that flooded the zone.
Figure 1: Spectral signature of a CRISM cube, the positive slope at 2000 nm is typical of chlorides.
Figure 2: Sample of the typical structure of the flat area where chlorides accumulate.
This distribution of chlorides differs from the typical emplacement of martian chlorides defined by [11], these authors showed that most of the chlorides, although being local depocenters, are accumulated in slopes and closed systems. The deposits in Icaria Planum are widespread and scattered without relation to the slope, which implies they are similar to playa like deposits, as is the case in the Uyuni salt flat [12].
The Uyuni salt flat is not only similar to the Icaria Planum deposits because of its formation, Uyuni is located high in the Andean altiplano, at 2670 m.a.s.l, which creates extreme conditions regarding radiation, temperature, precipitation, and salinity [12]. This resemblance makes Uyuni an integral analogue of at least some of the chloride-rich zones of Mars.
Conclusions
The mapping of the Icaria Planum deposits showed that these chlorides are similar to the Uyuni salt flat, which showcases the importance of the last one as a planetary analogue, and the diversity of salt formation processes that acts in the martian surface.
References
[1] McKnight et al. (2023). Distinct Hydrologic Pathways Regulate Perennial Surface Water Dynamics in a Hyperarid Basin.
[2] Cabrol, N. A., & Grin, E. A. (2003). Overview on the formation of paleolakes and ponds on Mars.
[3] Osterloo et al. (2008). Chloride-Bearing Materials in the Southern Highlands of Mars.
[4] Christensen et al. (2004). The Thermal Emission Imaging System (THEMIS) for the Mars 2001 Odyssey Mission.
[5] McEwen et al. (2007). Mars Reconnaissance Orbiter’s High Resolution Imaging Science Experiment (HiRISE).
[6] Murchie et al. (2007). Compact Reconnaissance Imaging Spectrometer for Mars (CRISM) on Mars Reconnaissance Orbiter (MRO).
[7 ]Roy et al. (2014). Landsat-8: Science and product vision for terrestrial global change research.
[8] Aileen Yingst et al. (2023). A Geologic Map of Vesta Produced Using a Hybrid Method for Incorporating Spectroscopic and Morphologic Data.
[9] Penasa, L., & Brandt, C. H. (2021). europlanet-gmap/mappy: Latest (latest) [Software].
[10] Tanaka et al. (2014). Geologic map of Mars.
[11] Leask, E. K., & Ehlmann, B. L. (2022). Evidence for Deposition of Chloride on Mars From Small‐Volume Surface Water Events Into the Late Hesperian‐Early Amazonian.
[12] Risacher, F., & Fritz, B. (1991). Quaternary geochemical evolution of the salars of Uyuni and Coipasa, Central Altiplano, Bolivia.
How to cite: Jimeno Ruiz, N. and Suarez Valencia, J. E.: Mapping of chloride deposits on Icaria Planum on Mars, and its possible correlation with the Uyuni salt flat in Bolivia, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-483, https://doi.org/10.5194/epsc2024-483, 2024.
The objective of this series of experiments was to support the data interpretation of short period changes in insolation of planetary bodies such as solar eclipses or transits. The temperature response of Martian regolith to the transit of Phobos has been observed by InSight and interpreted in terms of layering in the near surface [Mueller et al. 2021]. Current and future mission may observe similar transits if a suitable thermal infrared instrument is available. An opportunity to observe a similar event with the roles of Mars and Phobos switched may be the rover on the Martian Moons eXplorer (MMX) carrying the Mini-RAD radiometer [Ulamec et al. 2023]. The original mission plan did not foresee a landing on the Mars facing side, but in case the plan is revised due to the changing launch date there might be a chance. There are however open questions about the interpretation of such observations using 1D models of heat conduction [e.g. Mueller et al. 2021]. These models assume that the subsurface is a continuum, while in reality the material consists of particles that are not necessarily small compared to the characteristic depth scale (skin depth) of the material responding to the changes in insolation. The Planetary Ices Laboratory has the necessary equiment to recreate such events on Earth in form of a thermal vacuum chamber and a solar simulator (Fig. 1). The chamber walls can be cooled with liquid nitrogen and a window with mirror on the chamber ceiling allows illumination of an approximately 15 cm diameter spot by the solar simulator. For our test set-up we have placed a tray of about 38x18x2cm filled with regolith analogue material in the center of the chamber (Fig. 2) i.e. encompassing the illuminated spot. The radiometer is set up within the chamber on a rotary stage that allows to move the Field of View of the instrument along the long axis of the tray. The FoV is small enough to fall within the illuminated spot and the rotary stage allows adjustments within the running setup. A small TIR camera also observed the tray. This instrument has lower temperature resolution but provides data on temperature inhomogeneity of the sample. The general approach of simulating the eclipses was to first evacuate the chamber, either refill with 5 mbar CO2 or to leave at <1e-4 mbar, to cool the chamber to below -60 °C, then to illuminate the regolith analogue material until the temperature approaches an equilibrium. After that the solar simulator shutter is closed for a certain duration to simulate an eclipse. This is a more abrupt change in insolation than in reality, however the purpose of these tests is to explore the limits to which our 1-D heat conduction models, that are based on the continuum assumption, can reproduce the response of materials where the layer thicknesses are not large compared to the particle sizes. For this the duration of the insolation excursion is much more important than the shape of the excursion, because the duration governs the depth to which the temperature response extends. The simple on-off state of the insolation will simplify the model calculations. During our time at the planetary ices lab we conducted 8 different tests, consisting of 4 different material configurations and 2 different pressures, corresponding to Mars and space. The first configuration consists of 2-4 mm Mojave Mars Simulant (MMS) filling the tray to a depth of 2 cm. This is intended to correspond to the deeper layers of Mars Regolith. In reality the particle size distribution of regolith is wider but we chose the coarser material to achieve a greater contrast between the layers in the following. The next configuration distributes MMS sieved to < 100 µm over the coarse material. The volume was chosen to correspond to 2 – 3 monolayers at this particle size. This material was chosen to represent airborne Martian dust that is continually removed by dust devils (or, infrequently, by landing rockets) and resettles during the dusty season and dust storms. The next configuration adds about 3 mm more dust by extending the rim of the tray and filling the tray to the new height. This was chosen as an example of a thicker dust layer. Finally, there is a concern that larger clasts (~cm), which can be expected to fill a significant fraction of the FoV, may affect the data inversion, as they cannot be easily modeled using a 1D model. To this end we added such clasts on top of the dust layer. The eclipses followed a sequence of durations that (with exceptions) were 5 sec, 10 sec, 20 sec, 40 sec, 120 sec. For cases where the top layer thermal conductivity was very low we extended the last eclipse to 240 sec. After each eclipse the insolation was left on for at least 3 times the eclipse durations so that temperatures could return to near the pre-eclipse values. The temperature response of the different durations extends to different depths an therefore may sense different layers with different thermophysical properties. The surface temperature was observed using a space qualified radiometer using a thermopile with a spectral bandpass from 8-14 µm at a sampling frequency of 1 Hz. The initial data analysis shows the effect of layering, but for a more complete evaluation the exact test conditions (chamber and sample temperature, solar simulator power) will have to be included in 1 D heat conduction models that are work in progress. Fitting these data under these well known conditions (boundary conditions, material properties) will show how accurate 1D continuum models are for the interpretation of short term changes in insolation such as transits. Depending on the results we will re-evaluate the interpretation of the InSight data. Acknowledgements: This project has received funding from the European Union's Horizon 2020 research and innovation programme under grant agreement No 871149.
How to cite: Müller, N., Knollenberg, J., Kaufmann, E., Hagermann, A., and Grott, M.: Analogue simulation of the temperature response of Mars surfaces to Phobos transits, and vice-versa, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1007, https://doi.org/10.5194/epsc2024-1007, 2024.
MIRS (MMX InfraRed Spectrometer) is the imaging spectrometer (0.9-3.6 µm) [1] of the JAXA MMX (Martian Moon eXploration) mission [2]. The mission will be launched in the autumn 2026 to the Martian system, with an arrival around Mars planned in autumn 2027. The main objective of the mission is to study the two Martian moons, Phobos and Deimos, to collect samples of Phobos and bring them back to Earth in 2031. Another major objective of the mission [3] and the MIRS instrument [1] is to answer key science questions regarding the transport processes of dust and water in the Martian atmosphere [3], such as: how do local and regional dust storms form, grow and evolve? What is the diurnal behaviour of water ice clouds (formation, transport, dynamics)?
The MMX probe will be injected into a quasi-circular equatorial orbit around Mars at an altitude of about 6000 km from the surface. From this particular orbit, three different observation modes of MIRS are expected for Mars observations (see Figure 1) in addition to the limb mode: the so-called nominal mode whose purpose is to monitor up to low-medium latitudes, global mapping mode that covers most of the lighted Martian disk up to high latitudes (± 60°), and region of interest mode that provides temporal resolution (down to 15 minutes) above a chosen area. Each mode will be used to study the spatial and temporal variations of the aerosols (atmospheric dust, water and CO2 ice), and their fine diurnal variations. Indeed, the particular orbit of MMX (the second probe after Hope from EMM [4] to be in equatorial orbit) will allow MIRS to provide high-resolution spectral images in near-infrared and at very different local times, which will certainly contribute to answer to the questions addressed above. 
Figure 1: Illustration of the three expected observation modes of MIRS in addition to the limb mode: [A] nominal mode, [B] global mapping mode and [C] region of interest mode. Credits: CNES.
The retrieval pipeline for the MIRS observations is being prepared in coordination between the instrument team and the Mars Sub Science Team of the MMX mission. The two retrieval modules for trace gas and aerosols with nadir and limb observations are currently under development and will be validated by testing them on existing OMEGA/Mars Express data (e.g., [5], [6], [7]). In this study, focused on the aerosol retrievals, we use the DISORT (DIScrete-Ordinate-method Radiative Transfer) code [8, 9] through the pyRT_DISORT Python module [10] to simulate the expected radiance of the Martian atmosphere that MIRS will measure. First, we will present the parameter space exploration (estimated at up to 10 dimensions) of the radiative transfer model done to quantify the impact of each physical parameter (e.g., observation angles, surface albedo) on the generated spectra. Then, we will discuss the look-up table created by using different algorithms to explore the parameter space and their sampling in order to optimise the size of the look-up table (estimated at almost 75 million spectra), the computation time to create it, and to search in it. This table will be used to retrieve the aerosol properties in the flight data, which allows a faster retrieval than spectrum-to-spectrum inversion (given the large quantity of data that will be returned by MIRS). Finally, we will present MIRS images simulated with DISORT in real conditions during Phase 0 (orbit insertion and phasing), corresponding to the arrival of the MMX probe around Mars in 2027, for the different MIRS observation modes. These simulated data will be used to test the retrievals with the look-up table for future observations of Mars, get an understanding of the reachable performances, and help in defining the Phase 0 Mars set of observations.
Acknowledgments:
We thank the MMX JAXA teams for their efforts and CNES for the financial support and collaboration to build the MIRS instrument.
References:
[1] Barucci M. A. et al. (2021) Earth, Plan. and Space, 73, 211. [2] Kuramoto K. et al. (2022) Earth, Plan. and Space, 74, 12. [3] Ogohara K. et al. (2021) Earth, Plan. and Space, 74, 1. [4] Amiri H. E. S. et al. (2022) Space Sciences Reviews, 218:4. [5] Nakagawa H. et al. (2022) Seventh International Workshop on the Mars Atmosphere: Modelling and Observations, id3562. [6] Aoki S. et al. (2024), Japan Geoscience Union Meeting 2024. [7] Kazama A. et al. (2024), The Tenth International Conference on Mars. [8] Connour K. & Wolff M. (2023) GitHub repository, pyRT_DISORT. [9] Stammes K. et al. (1988) Applied Optics, 27, 2502-2509. [10] Stammes K. et al. (2017) Astrophys. Source Code Library, 1708.006.
How to cite: Leseigneur, Y., Fourgeaud, L., Gautier, T., Lasue, J., Stcherbinine, A., Bertrand, T., Sawyer, É., Jouquey, M., Théret, N., Kazama, A., Aoki, S., Doressoundiram, A., Nakagawa, H., and Barucci, A.: Preparing Martian Atmospheric Observations with MIRS, the MMX Imaging Spectrometer, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-456, https://doi.org/10.5194/epsc2024-456, 2024.
NOMAD [1] (Nadir and Occultation for MArs Discovery) is a multi-channel spectrometer onboard the ExoMars 2016 Trace Gas Orbiter (TGO), which began its observations in April 2018. Among other two (LNO and UVIS), the Solar Occultation (SO) channel covers the infrared (IR) spectrum from 2.3 to 4.3 µm (2320 to 4350 cm-1). Composed of an echelle grating in Litrow configuration, a total of 6 diffraction orders (with a typical width from 20 to 35 cm-1) are selected during each solar occultation using an Acousto-Optical Tunable Filter (AOTF) with a sample rate of about ~1 s, allowing a vertical resolution of typically 1 km. In order to optimize the information content that can be retrieved from the data, we analyzed Level 1 calibrated transmittances [2, 3] from four diffraction orders: 134 (3011-3035 cm−1), 136 (3056-3081 cm−1), 168 (3775-3805 cm−1) and 169 (3798-3828 cm−1). At the IAA-CSIC we developed processing tools specifically designed to handle and remove some systematics present in the NOMAD data, such as spectral shift and bending on the baseline of the spectra [4]. In addition, we performed an in-house characterization of the measurement noise using covariance matrices, identifying the true random component of the noise. This improvement allowed us to obtain homogeneous vertical profiles from the surface to an altitude about ~120 km. The profiles shown here have been retrieved combining two diffraction orders, this is, performing a global fit using spectra from two orders simultaneously. We combined pairs of orders 134 or 136 with 168 or 169 using them at different altitude ranges. The first two contain relatively weak absorption lines (S~10−21 cm−1/(molecule·cm−2) allowing the sample of the low atmosphere. On the contrary, the second set of orders contain strong lines close to center of the ν3 band(S~10−19 cm−1/(molecule·cm−2) which are useful to sample the upper atmosphere. This methodology is possible only when both orders have been measured during the same solar occultation, and although it limits the number of occultations available, it is necessary in order to avoid optically thick lines. Typically, we used low altitude orders (134/136) below 60 km and high altitude orders (168/169) above 60 km.
The content presented here is a follow-up work building upon several previous studies [5–8]. We extended the dataset selecting a total of 6561 occultations taken during Martian Years (MY) 34,35 and 36. We discuss detailed seasonal and latitudinal maps, showing the vertical distribution of the water vapor abundance and its variability thought the year. A summary of this work can be seen in Fig. 1, where we show the seasonal variation of all the retrieved water vapor profiles.
Figure 1: Vertical distribution of water vapor during MYs 34, 35 and 36 for the Northern (middle panel) and Southern (bottom panel) hemispheres. Horizontal axis shows the Solar longitude. Top panel indicates the latitude and local time of the observations.
Figure 1 clearly shows a repeated pattern in both hemispheres. Water vapor is present in a more vertically extended range during the Southern summer (perihelion season) whereas it is mostly confined to low altitudes (below 20 km) during the aphelion season, which corresponds to the Northern summer. Also, the characteristic Global Dust Storm (GDS) that took place in 2018 can be seen at the beginning of the MY 34 (LS~190º), showing a distinctive peak in water abundance in the northern hemisphere that is not repeated again.
In addition, we analyzed in detail the latitudinal distribution of water vapor during the perihelion season. We noticed a strong vertical plume at 60ºS - 50ºS injecting H2O into the mesosphere, reaching abundances of about ~50 ppmv at 100 km. We observed this event repeatedly in the three Martian years analyzed, although with inter-annual variations in both its magnitude and timing. The plume showed a weaker structure with less abundance during MY 34. We suggest that this difference respect to MYs 35 and 36 could possibly due to indirect effects of the MY 34 GDS. A summary of this analysis is presented in Fig. 2 where a storng water vapor injection can be seen in MYs 35 and 36 (panels B and C respectively).
Figure 2: Water vapor latitudinal variation during LS = 260º-280º for MYs 34, 35 and 36 (panels A, B and C respectively). Dots in top panels indicate latitude, Solar Longitude and Local Solar Time of the observations.
Acknowledgments
The IAA/CSIC team acknowledges financial support from the Severo Ochoa grant CEX2021-001131-S and by grants PID2022-137579NB-I00,RTI2018-100920-J-I00 and PID2022-141216NB-I00 all funded by MCIN/AEI/ 10.13039/501100011033. A. Brines acknowledges financial support from PRE2019-088355 funded by MCIN/AEI/10.13039/501100011033. ExoMars is a space mission of the European Space Agency (ESA) and Roscosmos. The NOMAD experiment is led by the Royal Belgian Institute for Space Aeronomy (IASB-BIRA), assisted by Co-PI teams from Spain (IAA-CSIC), Italy (INAF-IAPS), and the United Kingdom (Open University). This project acknowledges funding by the Belgian Science Policy Office (BELSPO), with the financial and contractual coordination by the ESA Prodex Office (PEA 4000103401, 4000121493), by Spanish Ministry of Science and Innovation (MCIU) and by European funds under grants PGC2018-101836-B-I00 and ESP2017-87143-R (MINECO/FEDER), as well as by UK Space Agency through grants ST/V002295/1, ST/V005332/1, ST/Y000234/1 and ST/X006549/1 and Italian Space Agency through grant 2018-2-HH.0. This project has received funding from the European Union’s Horizon 2020 grant No 101004052. US investigators were supported by the NASA. Canadian investigators were supported by the Canadian Space Agency. We want to thank M. Vals, F. Montmessin, F. Lefevre, F. Forget and the broad LMD/IPSL team supporting the Mars PCM.
References
1. Vandaele, A. C. (2018). Space Science Reviews, 214.
2. Thomas, I. R. (2022). Planetary and Space Science, 218.
3. Trompet, L. (2023). Journal of Geophysical Research: Planets, 128(3).
4. López-Valverde, M. A. (2023). Journal of Geophysical Research: Planets, 128(2).
5. Brines, A. (2023). Journal of Geophysical Research: Planets, 128(11).
6. Aoki, S. (2019). Journal of Geophysical Research: Planets, 124(12).
7. Aoki, S. (2022). Journal of Geophysical Research: Planets, 127(9).
8. Villanueva, G. L. (2022). Geophysical Research Letters, 49(12).
How to cite: Brines, A., Lopez-Valverde, M. A., Stolzenbach, A., Modak, A., Funke, B., González-Galindo, F., López-Moreno, J. J., Sanz-Mesa, R., Aoki, S., Vandaele, A. C., Daerden, F., Thomas, I., Erwin, J., Trompet, L., Ristic, B., Villanueva, G., Liuzzi, G., Patel, M., and Bellucci, G.: Global Vertical Distribution of Water Vapor in the Martian Atmosphere for 6 years of ExoMars-TGO/NOMAD observations, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1018, https://doi.org/10.5194/epsc2024-1018, 2024.
To inform future Mars spacecraft development, it is important to understand the impact of new observational platforms on our ability to measure Mars’ global climate properties and to identify weather patterns in the Martian atmosphere. In particular, recent studies have begun to focus on constellations of multiple orbital weather-monitoring satellites in anticipation of new missions to Mars in the coming decade [1-4].
We present a series of Observing System Simulation Experiments (OSSEs) that were carried out to determine the optimal configuration and payload for potential future weather-monitoring satellites. This is based on the ability of a data assimilation system to reproduce the atmospheric state under clear, global dust storm, and regional dust storm conditions, while accounting for observational uncertainties.
Method
To simulate the impact of atmospheric observations that do not yet exist for Mars, we used an “Observing System Simulation Experiment” (OSSE) framework. This approach has been applied for more than 30 years to evaluate the effects of future instruments on monitoring the Earth’s climate system. Our OSSE framework had three components:
First, the Nature Run (NR) was a GCM simulation at higher-than-standard spatial resolution, which represented the true state of the atmosphere. The point of using the OSSE framework is that the true state of the atmosphere is known exactly; something that is not possible with real atmospheres. This allows us to isolate the differences between observing platforms rather than focusing on the differences between model and reality. For the NR we used the Mars PCM [5] at 2o x 2o horizontal resolution. The NR covered three meteorologically different 60o-Ls periods during Mars Year (MY) 34, shown in Fig. 1: a clear period with no dust storms (Ls = 60-120o), a period with a global dust storm (Ls = 180-240o), and a period with a regional dust storm (Ls = 300-360o).
Figure 1: Mars Year 34 infrared column dust optical depth as a function of time of year and latitude [6]. The three study periods are indicated by boxes [4].
Second, the synthetic observations were created to mimic as closely as possible what an instrument would observe if it were observing the simulated Mars atmosphere. Our synthetic observations represented retrievals of atmospheric temperature profiles, dust opacity profiles, and column dust optical depth measurements made by the existing instruments MRO-MCS, ExoMars TGO-ACS-TIRVIM, and EMM-EMIRS, and proposed or desirable future instruments, such as a low-altitude near-polar orbiter that cycles through the full range of local times over 14 sols, a 5720-km altitude low-inclination orbiter, and areostationary orbiters, along with various combinations of these spacecraft. The experiments varied the instruments studied (point and imaging radiometers observing in the thermal infrared) and observational strategies (nadir, off-nadir, and limb).
Finally, the data assimilation scheme combined the observations and GCM at “standard” spatial resolution (5.625o x 3.75o) into a sequence of “reanalyses”, and we compared these reanalyses with the “true” state of the atmosphere in the NR. We used the Local Ensemble Transform Kalman Filter (LETKF) with 36 ensemble members [7]. Figure 2 illustrates the process schematically.
Figure 2: Flow diagram showing an Observing System Simulation Experiment for Mars [4].
Results
We quantified the accuracy of each reanalysis by computing several diagnostics, foremost among which was the root-mean-square difference (RMSD) between the reanalysis and the NR (“truth”). This was computed for observed variables (temperature, dust) and for unobserved variables (surface pressure, horizontal wind components). The latter are updated dynamically by the GCM during the assimilation process so that they converge (in principle) to a state where they are in balance with the measured temperature field, within observational uncertainty. We also studied higher order diagnostics such as the amplitude and phase of the thermal tides.
Figure 3 shows an example of the column dust optical depth RMSD throughout the global dust storm period, for several instrument configurations.
Figure 3: Visible column dust optical depth RMSD normalised to 610 Pa during the global dust storm period of Mars Year 34 [4]. The RMSD is an area-weighted, global average of the difference between the reanalysis and the Nature Run. Individual lines include a GCM simulation using a climate dust scenario with no observations assimilated (black), MRO-like observations (magenta), a near-polar low altitude orbiter (grey), three areostationary platforms (olive), and combinations of areostationary and other orbiters (yellows and greens). Some other lines show configurations that will not be included in the presentation. The dashed black line is the observational error.
In dust storm conditions, the ability of the assimilation to correctly reconstruct the weather generally increases with the comprehensiveness of the observing constellation, in terms of the number of platforms, spatial coverage, and variety of observations. While several orbital configurations perform equally well when looking at global or daily-averaged meteorological fields, significant differences between configurations emerge when looking at the diurnal cycle or at regional scale. Large observational uncertainties greatly affect the accuracy of the results.
Outlook
This approach can be applied to constellations of SmallSats focused on Mars weather monitoring, assuming the availability of simple payloads dedicated to temperature and dust observations. The assessment of the effect of direct wind measurements (e.g., using a Doppler LIDAR or sub-mm sounder on larger orbital platforms) and the impact on the weather forecasting capability would be valuable topics for follow-on studies.
Acknowledgments
This work was carried out partly under the ESA SWIMALONG project, contract number 4000138021/22/NL/DB. RY was supported by UAE University grants G00003322 and G00003407.
References
[1] Montabone et al. (2020), NASA White Paper, Bull. AAS, 53, 281, 10.3847/25c2cfeb.0cdca220.
[2] Montabone et al. (2020), Access2Space Workshop, JHU APL.
[3] Montabone et al. (2021), EPSC2021-625, 10.5194/epsc2021-625.
[4] Montabone et al. (2023), ESA report SWIM-OU-TN-002, contract no. 4000138021/22/NL/DB.
[5] Forget et al. (1999), JGR Planets, 104(E10), 24155-24175, 10.1029/1999JE001025.
[6] Montabone et al. (2020), JGR Planets, 125, e2019JE006111, 10.1029/2019JE006111.
[7] Young et al. (2022), JGR Planets, 127, e2022JE007312, 10.1029/2022JE007312.
How to cite: Young, R., Montabone, L., Millour, E., Paardekooper, D., Parfitt, C., and Wilson, C.: Assessing the impact of present and future orbiters on Mars weather monitoring using data assimilation, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1029, https://doi.org/10.5194/epsc2024-1029, 2024.
Introduction: Martian dust storms play a crucial role in the red planet's climate and weather patterns, affecting both atmosphere and surface conditions over various time scales [1]. Orbital data help track these storms, which pose challenges for spacecraft operations and energy production [2], [3], [4] and [5]. Although monitoring is precursor, accurate forecasting is essential for future Mars exploration [6]. Supervised machine learning (SML) has already been used for automatic detection of Martian dust storms from visible images [7], [8]. However, we are not aware of SML and/or unsupervised machine learning (UML) being utilized for this application with other data sources.
Data: Montabone et al. (2015) have already identified a certain degree of interannual repeatability of large-scale dust storms from zonal mean of column dust optical depth normalized to the reference 610 Pa pressure level (CDOD@610) [9]. A novel approach relying exclusively on UML (see Figure 1) has been developed to detect, aggregate and track Martian large-scale dust events both in space and time (ST-DATMADE), from the publicly-available, multiannual (from Martian Year (MY) 24 to 36) CDOD@610 dataset described in Montabone et al. (2015, 2020) [9], [10]. Large-scale dust events are characterized by surface areas ≥1.6×106 km2 and last more than two Martian days [2]. Normalized CDOD is used to remove topographic features (plains, volcanoes etc.) which affect the distribution of dust within the column. The dataset contains gridded observation maps and kriged maps from CDOD satellite retrievals. Being spatially interpolated, regularly kriged maps (see Figure 2, 1st column) are more relevant to use for SML application compared to irregularly gridded maps.
Detection: Detection of dust events involves the identification of spatio-temporal anomalies in CDOD maps. A dust event is defined as a temporal series of dust episodes characterized by foreground dust anomalies. A foreground dust anomaly refers to a significant increase in CDOD observed during a given Sol (Martian day), regardless of spatial dimension, relative to the background dust level. Background dust refers to the typical or baseline level of diffused dust present in the Martian atmosphere. Therefore, to facilitate the recognition of these anomalies, a “background/foreground” segmentation is initially performed to enhance subsequent data partitioning (see Figure 1). The CDOD data distribution is adjusted by mapping the lowest data associated with background dust into a threshold value, , calculated as follows:
The low dust loading season (LDLS) is a multiannual period between LS (Solar Longitude) = 10° and LS = 140° where there is almost no large-scale dust injection [11]. Here, the LDLS background, , is defined as the 95th percentile () of CDOD during LDLS between MY24 and MY36. It represents the multi-annual atmospheric background during the “clear” season. Additionally, the daily background, , is defined as the 33rd percentile () of CDOD at a given Sol and MY. It considers that the background dust may evolve as dust events may increase the surrounding dust opacity.
Furthermore, the adjusted dataset distribution is extremely and positively skewed, in such case an inverse transform may help to reduce the skewness, making the distribution more symmetrical and evenly distributed, particularly when addressing extreme dust events, [12]. This brings extreme values closer to average values. Within an episode of foreground dust anomaly, distinct features are identified as dust core and dust cloud. A dust core is a value-based partition where atmospheric dust injection is likely. A dust cloud is a value-based partition where atmospheric dust diffusion is likely. Grouping between background dust, dust cloud and dust core is realized through CDOD values partitioning in k=3 partitions using a widely-known UML algorithm: k-means [13]. See Figure 1 and Figure 2, 2nd column.
Aggregation: Density-Based Spatial Clustering of Applications with Noise (DBSCAN) algorithm ([14]) is used for spatial clustering of the partition with the highest CDOD values (i.e. dust core). A dust instance is an individual, isolated spatial-based cluster within the dust core partition, produced by DBSCAN (see Figure 1 and Figure 2, 3rd column).
Tracking: A series of space-coherent and time-continuous dust instances is organized into a dust sequence. To track instances from one Sol to another (Sol-to-Sol) the pairwise centroid distances of consecutive Sols is compared to a “sphere of membership” (SOM). As well as the percentage of overlapping (OLP) of instances from consecutive Sols. Thereby, it is possible to track dust event instances Sol-to-Sol and to deduce whether they form sequences (see Figure 1 and Figure 2, 4th column).
Catalog: Previous section of tracking provides Sol-to-Sol cluster assignments (instances organized in sequences) allowing to build a catalog of dust event instances tagged with an “id” with format: MY.._Sol_..._1,2,3 etc. Space-consistent and time-continuous instances are labelled with a sequence “id” as: MY.._A,B,C, etc. (see Table 1).
Statistics: This catalog will allow a comprehensive analysis of large-scale dust events based on CDOD data. In particular, their spatio-temporal distribution through trajectories, region of origination, LS of origination etc. as well as intensity distribution through area and CDOD statistical features (mean, median etc.).
Acknowledgments: The authors would like to acknowledge the use of the publicly available dataset on the LMD webpage: http://www-mars.lmd.jussieu.fr/mars/dust_climatology/index.html and the NASA PDS webpage:
https://pdsatmospheres.nmsu.edu/data_and_services/atmospheres_data/MARS/montabone.html.
References:
[1] L. Montabone and F. Forget, 2018, http://hdl.handle.net/2346/74226
[2] B. A. Cantor et al., 2001, https://doi.org/10.1029/2000JE001310
[3] H. Wang and M. I. Richardson, 2015, https://doi.org/10.1016/j.icarus.2013.10.033
[4] M. Battalio and H. Wang, 2021, https://doi.org/10.1016/j.icarus.2020.114059
[5] C. Gebhardt et al., 2022,
https://www-mars.lmd.jussieu.fr/paris2022/abstracts/poster_Gebhardt_Claus.pdf
[6] L. Montabone et al., 2022,
https://www-mars.lmd.jussieu.fr/paris2022/abstracts/oral_Montabone_Luca.pdf
[7] R. Alshehhi and C. Gebhardt, 2022, https://doi.org/10.1186/s40645-021-00464-1
[8] K. Ogohara and R. Gichu, 2022, https://doi.org/10.1016/j.cageo.2022.105043
[9] L. Montabone et al., 2015, https://doi.org/10.1016/j.icarus.2014.12.034
[10] L. Montabone et al., 2020, https://doi.org/10.1029/2019JE006111
[11] F. Forget and L. Montabone, 2017, http://hdl.handle.net/2346/72982
[12] B. G. Tabachnick and L. S. Fidell, Using Multivariate Statistics. Pearson Education, 2013.
[13] S. Lloyd, 1982, https://doi.org/10.1109/TIT.1982.1056489
[14] M. Ester et al., 1996, https://www2.cs.uh.edu/~ceick/7363/Papers/dbscan.pdf
How to cite: Lombard, T. and Montabone, L.: Spatio-Temporal Detection, Aggregation and Tracking of Martian Large-Scale Dust Events, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1334, https://doi.org/10.5194/epsc2024-1334, 2024.
Introduction
Cyclic absorption of solar radiation generates oscillations in atmospheric fields. These oscillations are called atmospheric or thermal tides, which are furthermore modified by topography, surface properties, and atmospheric absorber concentration. This leads to a complex mix of sun-synchronous and non sun-synchronous tides that propagate around the planet eastward and westward. A major aspect of the study is to obtain some global insight into the Martian atmospheric ter-diurnal oscillations seen at a very small collection of lander sites using a Mars Climate Database (MCD) General Circulation Model (GCM).
Observations and Methods
Here we used hourly binned surface pressure observations by Mars Science Laboratory (MSL), InSight, Viking Lander (VL) 1, and VL2 to calculate harmonic components with a Fast Fourier Transform (FFT) for each location. A window of three sols was used to get at least one observation for each hour. The binned value was averaged in case of multiple observations per hour. Amplitudes and phases were then calculated for the middle sol of the three-sol window. The analysis was performed for each station by sliding the window over all the pressure data that was available for each platform.
For the MCD GCM surface pressure data, we used a similar approach except that one-sol window was used for each location. These global GCM fields allow for mapping of the evolving tide harmonic results and, further, decomposing these into constituent eastward and westward propagating components. That was done by taking a 2D FFT (in time and longitudes) for the diurnal pressure fields. Moreover, we used special GCM simulations with and without the effect of radiative heating by water ice clouds provided by the LMD team.
Results
MSL and InSight showed very similar seasonal cycles with highest ter-diurnal amplitudes at about Ls 60◦ , Ls 130◦ , and Ls 320◦ (Figure 1) . The latest amplitude peak corresponds to a ”C” dust storm and is also visible in the diurnal and semi-diurnal tidal amplitudes. However, during the ”A” dust storm (at about Ls 230◦) neither platforms detected amplitude spikes. Observed VL1 and VL2 ter-diurnal amplitudes seem to be lacking a clear structure, but VL1 showed quite repeatable and similar pattern as MSL and InSight during the first half of the year. In addition, global dust storms during MY 12 (1977b) and 15 were clearly detected by the ter-diurnal tide amplitude. However, ter-diurnal amplitude did not show a clear response during the 1977a planet-encircling dust storm. The global dust storm during MY 34 was visible in the MSL pressure data demonstrating the highest ter-diurnal amplitude during the MSL mission by now.
Figure 1: Ter-diurnal amplitudes (normalized by the diurnal mean pressure, %) from the MSL, InSight, VL1, and VL2 observations as well as from MCD as a function of the season (Ls).
MSL and InSight showed very similar pattern in ter-diurnal phases as well, with values between 05-07 LTST most of the year and 08-10 LTST during Ls 60◦-120◦ (Figure 2). Phases observed by the Viking Landers were lacking a clear structure, but VL1 showed some pattern at the start of the year. The phases predicted by the MCD data were in quite good agreement with observations, except for the MSL.
Figure 2: Same as Figure 1 but for the ter-diurnal phases (LTST).
Overall, calculations on MCD pressure data seem to underestimate ter-diurnal amplitudes throughout the year, especially during the dusty season. Globally, MCD predicted weakest amplitudes at the equinoxes, while strongest ones were predicted in summertime for both hemispheres with a clear wavenumber 6 pattern in longitude during northern hemisphere summer (Ls 90◦, Figure 3). During this time, model results suggest that the two most prominent modes are the sun-synchronous ter-diurnal tide (TW3) and an eastward propagating TE3 tide. The meridional structure of TE3 (Figure 4) mode is characteristic of a resonantly-enhanced Kelvin wave, with a uniform latitude structure.
Figure 3: Global ter-diurnal tide normalized amplitudes from the Mars Climate Database for Ls 0◦, Ls 90◦, Ls 180◦, and Ls 270◦.
Figure 4: Ter-diurnal amplitudes for westward (TW) and eastward (TE) propagating zonal wavenumber 1–4 components as a function of latitude from the MCD for each season.
Additional GCM simulations with and without radiative active water ice clouds suggest that TE3 mode amplitude peaks at the equator and is much less prominent without radiatively active water ice clouds (Figure 5). Moreover, the observed delay and advance of the ter-diurnal phase (Figure 2), as well as the amplitude patterns (Figure 1) indicate that the TE3 mode has a characteristic of the resonantly-enhanced ter-diurnal Kelvin wave. This is because MSL/InSight and VL1 are separated by about 180 degrees of longitude, and therefore the phase deviation due to the Kelvin wave would be similar at the 3 locations, as seen in the data. Interaction between solar radiation and zonal wave 6 variations in topography and surface properties (thermal inertia and albedo) may generate this mode. By contrast, TW3 tide was stronger away from the equator and was present for a much longer fraction of the year (Figures 4 and 5). This component may be a result of direct solar forcing or nonlinear interaction between diurnal and semi-diurnal tides.
Figure 5: Seasonal normalized amplitude evolution of TW3 and TE3 modes with (left panel) and without (right panel) the effect of radiative heating by water ice clouds from the LMD GCM simulation. Dust column opacity contours are shown on the TW3 field and water ice cloud contours (limited to tropics) are shown on the TE3 field.
Based on this investigation, it seems to be reasonable to infer that ter-diurnal tide has a seasonal cycle at least near the equator, which is strongly influenced by water ice clouds. Nonetheless, the amplitude is typically lower than that for the diurnal and semi-diurnal tides. Definitely more modeling studies are needed to understand the behavior and forcing mechanisms of this important component of the Martian atmospheric tide.
Reference
Joonas Leino, Ari-Matti Harri, John Wilson, et al. Ter-diurnal Atmospheric Tide on Mars. ESS Open Archive . April 26, 2024. DOI: 10.22541/essoar.171415900.05448468/v1
How to cite: Leino, J., Harri, A.-M., Wilson, J., Banfield, D., Lemmon, M., Paton, M., Rodríguez-Manfredi, J., and Savijärvi, H.: Ter-diurnal Thermal Tide on Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-36, https://doi.org/10.5194/epsc2024-36, 2024.