Introduction:  The response of a planet to loading is intimately linked to its interior structure. On Earth, studying the time-variable response of the lithosphere to the growth and decay of large ice sheets, or glacial isostatic adjustment (GIA), has been the standard approach to constrain the mantle structure [1]. However, such models have rarely been applied to the other terrestrial planets because of a lack of observational data [2].

Mars harbors two geologically young (<100 Ma) and large (~1000 km across) polar ice caps [3], which represent the only million-year-old surface features that can induce measurable surface deformations. Radar sounders have mapped these ice caps [4,5] and the lack of lithospheric flexure beneath the north polar cap (Fig. 1a,b) was used to constrain the interior of the planet to be cold with a stiff elastic lithosphere [4,5,7]. While the results of these studies are used to constrain the present-day thermal state of Mars [8], geodynamic models [9] struggle to explain the stiff lithosphere inferred at the north pole and, at the same time, young volcanism [10] and ongoing plume activity [11] (Fig. 1c).

Importantly, previous analyses have assumed the polar deformation to be at equilibrium, which is only valid if the time elapsed since the ice cap formed is greater than the time required for viscous adjustments. Previous calculations suggested that viscoelastic relaxation could allow thinner lithospheres and higher heat flows beneath the northern cap [4,5]. However, these estimations were based on simplistic constant mantle viscosity interiors, and only examined a single wavelength to represent the load, thereby providing limited insights into the potential effects of viscoelastic relaxation. Here, we investigate GIA on Mars in light of newly acquired InSight constraints on the planet’s interior structure [8–12]. Viscoelastic simulations of the polar deformations are coupled with thermal evolution models that account for InSight seismic measurements and observational constraints on recent volcanic activity [9–11].

Methods:  We have run a large suite of 3-D thermal evolution models using interior structure parameters constrained by InSight [9]. We test crustal thickness models [12] and vary other key parameters such as the distribution of radioactive heat sources [13]. The occurrence of present-day melting in these simulations is used to select acceptable models. In total, 17 models with a wide range of crustal structures are selected (Fig. 1). None of these models match the results from [4,5], and in absence of external constraints, all these models are equally possible.

To explore the extent of relaxation beneath the north polar cap, Love numbers are computed using ALMA [14]. Therein, the momentum equation of an incompressible Maxwell viscoelastic planet is solved applying a propagator method. Love numbers are calculated up to Legendre degree 50, which is sufficient to resolve the polar cap loading (~degree 8, ~1250 km wavelength).

For each of the 17 models, the predicted internal structure beneath the northern cap (>60°N) is discretized into 68 constant-thickness layers with given viscosity, density, and rigidity, from the core to the surface (Fig. 2). The time evolution of the deformation for these models is predicted using ALMA and the polar cap load potential [5,7]. Our model assumes that the ice cap formed instantly at a given time in Martian history. We also estimate the rate of deformation and the associated time-variable degree-2 zonal gravity coefficient (J₂). A model is deemed acceptable if the predicted viscoelastic deformation matches radar observations [5].

Results and conclusions:  The present-day interior structure beneath the northern cap shows mantle viscosities ranging from 10²¹ to 10²⁴ Pa s (Fig. 2). Given the large interior viscosity, the equilibrium response of the interior to loading is achieved over long-timescales (>100 Myr). The north polar cap is young (<10 Ma, [15]), which implies that GIA is presently occurring on Mars. As a result, the elastic lithosphere thickness estimated in [4,5] are overestimated. In particular, some geodynamic models that were deemed unsuccessful because of their mismatch to the north polar elastic thickness are now acceptable (Fig. 1-2).

Only a subset of models predicts small polar deformations consistent with radar observations [4,5], and these are characterized by a thick crust (≥50 km on average), cold lithosphere (elastic thickness >150 km), and large mantle viscosity (>10²² Pa s). A thick crust is consistent with the independent reassessment of InSight seismic data [16]. Models with a higher crustal heat production cool faster, have higher mantle viscosities, slower interior response to loading and thus better fit radar observations. For accepted models, GIA in the polar regions show rates >10^-3 mm/year and strain rates >10^-20 s^-1.

Our analysis further constrains the bulk of the northern cap to have formed no more than 7 Myr ago, which is in remarkable agreement with independent Global Climate Model [15]. Models with a linear ice cap formation would imply even younger ages. If the mantle viscosity were precisely known, our framework could be used to place constraints on the formation history of the ice cap.

GIA affects the long-wavelength gravity field. Over the last 20 years, J₂ may have varied by up to 10^-11 m s^-2. Together, our results implies that GIA could be measured from precise laser ranging or with a dedicated GRAIL-like gravity mission to Mars [17,18]. Alternatively, reanalyses and extended collection of radio tracking data could enable discerning a trend in the gravity field.

References: [1] Peltier W. R. (1974). [2] Greve R., et al. (2003) [3] Byrne S. (2009). [4] Phillips R. et al. (2008). [5] Broquet A. et al. (2020). [6] Broquet A. (2020) AB-Ares/Te_HF_Conversion [7] Broquet A. et al. (2021) JGR:Planets. [8] Khan A. et al. (2021). [9] Plesa A.-C. et al. (2022). [10] Voigt J. R. C. et al. (2023). [11] Broquet A. & Andrews-Hanna J. C. (2023). [12] Wieczorek M. A. et al. (2022). [13] Hahn B. C. et al. (2011). [14] Spada G. (2008). [15] Levrard B. et al. (2007). [16] Kim D. et al. (2023). [17] Genova A. (2020). [18] Wörner L. et al. (2023).

How to cite: Broquet, A., Plesa, A.-C., Breuer, D., Wieczorek, M. A., Root, B. C., Klemann, V., Genova, A., and Andrews-Hanna, J. C.: Glacial isostatic adjustment on Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-847, https://doi.org/10.5194/epsc2024-847, 2024.

NASA’s InSight mission has brought new information about the Martian lithosphere (Banerdt et al. 2020), which warrants a re-analysis of the support of the crustal and sub-crustal masses. Furthermore, the discovery of a possible mantle plume underneath the region south of Elysium Mons (Broquet et al. 2023), gives evidence for recent magmatic activity underneath the crust of Mars causing dynamic support to volcanic structures.

The goal of this study is to combine a global flexural model combined with mantle convection modelling to study the different uncertainties in the geophysical parameters that dictate crustal (short-scale) and mantle (large-scale) mass anomalies. After conducting spectral analysis on the gravitational signal of we found that the Martian lithosphere can be best modelled by an elastic thin shell with the following physical parameters: crustal thickness of 60 ±10 km, crustal density of 3050 ±50 kg/m³, mantle density of 3550 ±100 kg/m³, and the elastic thickness (T_e) is found to be 80 ±5 km.

Figure 1: Crustal density from short scale gravity anomalies inversion. The flexural model is not able to represent to perfect isostasy ate Hellas Basin (and assumed present also at Utopia basin). Therefore, high crustal density is seen reflecting a shallower crust-mantle interface. Isidis and Argyle basin also show high mass regions that could be interpreted as shallow crust-mantle boundary or magmatic intrusions in the crust. Other high mass features are related to volcanoes and large impact craters. Finally, the northern polar planes show also anomalous high mass regions.

By using this spectral analysis, we were able to isolate the remaining long-wavelength signal that cannot be realistically modelled by a flexure model. The location of the residual anomaly correlates with the Tharsis Rise, which suggests active large-scale dynamic support of the volcanic region. A negative mass anomaly in the mantle underneath the Tharsis Region explains the remaining gravity residual. This anomaly could be interpreted as underplating of the lithosphere, a phase transition anomaly at 1000 km depth, or an rising mantle plume. Could mantle convection is still be active in Mars explaining the relatively young geologic surface volcanism?

The remaining short scale gravity residuals give insight in Martian crustal density distributions (Figure 1). Flexure models cannot account for these structures as the lithosphere is strong enough to reduce/negate any topographic signature. Especially in the northern polar plains several buried mass anomalies have been detected (Figure 2). The nature of these anomalies is unclear, as they could be interpreted with a volcanic origin, impact related structures, or tectonic orogeny that would all be buried be the sedimentary layer that is observed on the surface of the northern hemisphere. If these structures are interpreted as impact craters this would suggest of an older crustal age of the northern hemisphere of Mars than is now considered. New gravity satellite missions towards Mars are needed to uncover the nature of these subsurface mass anomalies (Wörner et al 2023, Genova et al 2020).

Figure 2: Contours  of the high mass anomalies in the nortern polar plains. These high mass anomalies seem to have no topograohic signature. Topography is shown in the colorscale and taken from the MOLA dataset.

References:

• Banerdt, W. B., Smrekar, S. E., Banfield, D., Giardini, D., Golombek, M., Johnson, C. L., … Wieczorek, M. (2020, March). Initial results from the InSight mission on Mars. Nature Geoscience, 13 (3), 183–189.
• Broquet, A., & Andrews-Hanna, J. C. (2023). Geophysical evidence for an active mantle plume underneath Elysium Planitia on Mars. Nature Astronomy, 7 (2),160–169.
• Genova, A. (2020), ORACLE: A mission concept to study Mars’ climate, surface and interior, Acta Astronautica, 166, 317-329.
• Wörner, L., B.C. Root, P. Bouyer, C. Braxmaier, D. Dirkx, J. Encarnação, E. Hauber, H. Hussmann, Ö. Karatekin, A. Koch, L. Kumanchik, F. Migliaccio, M. Reguzzoni, B. Ritter, M. Schilling, C. Schubert, C. Thieulot, W.v. Klitzing, O. Witasse (2023), MaQuIs—Concept for a Mars Quantum Gravity Mission, Planetary and Space Science, 239, 105800.

How to cite: Root, B., Alkahal, R., Qin, W., and Thieulot, C.: Exploration of high mass subsurface structures in the northern hemisphere with joint flexure and mantle convection modelling of the Martian gravity field , Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-730, https://doi.org/10.5194/epsc2024-730, 2024.

Impact basins, the largest type of impact craters, form during the early stages of planetary evolution in the first several hundred million years following planetary accretion. Their size and morphology are mainly controlled by impact speed, impactor mass and geological properties of the target at the time of impact. Significant progress has been made in comprehending impact basins in recent years, with most efforts aimed at the Moon, where basins are best preserved. However, our comprehension of Mars basins remains comparatively limited. To fill this gap, we used numerical simulations to study the formation of impact basins on Mars at different geological epochs and regions.

We simulated the formation of impact basins using iSALE-2D, a multi-material, multi-rheology shock physics code that can simulate impacts in geological materials [www.isale-code.de]. We considered vertical impacts of dunite impactors travelling at 10 km/s on a target composed of a basaltic crust and a dunitic mantle with a surface temperature of 213 K. Projectile sizes varied from 20 to 400 km in diameter. The target was considered flat for simulations involving impactors smaller than 200 km; larger impacts were evaluated using both flat and spherical target geometry. We adopted three temperature profiles representing the planet's thermal state at 4.4 Ga (soon after the crust formed, when it was relatively hot), 4.0 Ga (an intermediate state) and 3.5 Ga (when the crust was comparatively colder). Two crustal thicknesses were considered, 47 and 91 km, corresponding to the average values of the martian lowlands and highlands.

Our simulations were presented as cross-sections of half craters showing two distinct layers, crust and mantle. We enabled material tracing to follow the material movement. We measured the spacing of rings indicated by tracer displacements and compared their position to the overall morphology of the simulated basins. We observed the formation of a listric normal fault at the main rim, crosscutting the crust and mantle to depths comparable to the maximum excavation depth. This structure, referred to as "main fault", was observed across all basin sizes considered in this study, and thus was considered as a suitable marker for the final basin size. In sufficiently large basins, we observed an outer ring defined by a planar normal fault crosscutting crust and mantle to comparatively larger depths. This structure, referred to as "outer ring fault", was associated with a topographical high and a zone of localised crustal thinning, forming crustal blocks similar to the ones observed in regions of extensional tectonics.

The spacing of rings in our simulated basins is similar to the reported spacing of rings in lunar basins. On average, the inner and outer rings are positioned at  0.46 and 1.41 times the final basin radius respectively (Figure 1). This spacing is controlled by age and crustal thickness. Rings of basins formed in older basins or hotter crust tend to be positioned closer to the impact point than those formed in younger basins or colder crust. Crustal thickness excerpt a similar effect, especially on the position of outer rings. Outer rings of basins formed in a thin crust are ~14% closer to the impact point than outer rings of basins formed in a thick crust.

Our results have implications for the formation and evolution of impact basins on Mars, helping to deepen our understanding of basin’s inner structure and multi-ring basin formation. Moreover, our results can be used to constrain the thermal state and crustal structure on Mars at the time of basin formation by comparing our findings with the observed morphology of Martian basins. This provides important insights to the bombardment history of Mars and to our knowledge of large impacts in the early Solar System.

Figure 1 - Relative ring spacing of the inner (r_inn) and outer (r_out) rings measured from 60 numerical simulations of impact basins on Mars created by projectiles from 20 to 200 km in diameter. The plots present, from top to bottom, the complete dataset, and simulations for 3.5, 4.0 and 4.4 Ga. Inner ring ratios for a thin and thick crust are presented in orange and blue respectively; outer ring ratios for a thin and thick crust are presented in red and green respectively. In the rare cases more than one outer ring was present, only the inner most radii were recorded.

How to cite: Branco, H., Miljkovic, K., and Plesa, A.-C.: Ring formation in impact basins on Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-288, https://doi.org/10.5194/epsc2024-288, 2024.

Introduction: Since the 60’s, extensive surveys with orbiters have produced large quantity of data on planet Mars, including very high resolution products (e.g. [1-2]) that allow precise studies of the surface geomorphology. Quantitative approaches such as morphometry have been broadly used on orbital imagery to produce impact craters scaling laws [3-5]. These laws are needed to quantify erosion (e.g. [6]), and by extension to better understand the surface processes that occur on the red planet. However, only one study has focused on crater diameters below 1 km [5] and scaling laws for diameters below 50 m have not been investigated yet. In our study, we counted the craters and measured their depths and diameters on landslide deposits located in Valles Marineris [7]. We produced a scaling law of depth d as a function of diameter D for craters below 50 m.

Methods: We used High Resolution Imaging Science Experiment (HiRISE, [1]) products to count all the craters in a specific area. We focused on ESP_024255_1750, a HiRISE image and the associated Digital Elevation Models (DEM, [2]), which both have a resolution of 0.5 m/pixel. We mapped the landslide deposits, which are restricted to the northern part of the image, at the base of the canyon walls. The data have been studied with Geographical Information System (GIS) Arcmap 10.8. We first counted all the craters at the surface and made diameters measurement using the HiRISE image. Then, depths have been measured following our handmade model through the ModelBuilder tool. The model fitted a local surface for each crater from elevations encompassing the crater rims, on a 3 m buffer area. We finally subtracted the elevations from the local surface and the DEM to obtain the crater depth.

Results and discussion: We counted 32128 craters of all diameters on ESP_024255_01750 within the landslide region. We reduced this count to keep all the craters with D > 10 m and excluded the negative depths that are due to DEM vertical uncertainties. This yielded a final count of 1252 craters. We fitted the freshest craters to get the morphometry under the form d = cD^α following the methods from Boyce and Garbeil, 2007 [3] (figure 1). This produced the following expression:

d = [0.137±0.290] D^{[0.862±0.419]} ,

where d and D are depths and diameters in meters. Parameter α is slightly lower than the value found in Watters et al., 2015 (α = 0.205±0.012, [5]), highlighting that craters are shallower in this range of diameters. Parameter c has an intermediate value between those from Watters et al., 2015 (c = 1.012±0.009 [5]) and Garvin et al., 2003 (c = 0.81 [4]). These results have to be compared with our previous results in which we found a higher value for c (c = 1.014, Millot et al., 2024, conference paper submitted to the 10^th International Conference on Mars [8]) on Interior Layered Deposits (ILDs, e.g. [9]), light-toned sedimentary rocks that are known to be easily erodible [6]. These differences highlight the impact of terrains on the scaling laws. Finally, parameters c and α have high uncertainties due to the small area involved in the count: we expect these uncertainties to decrease if we include more landslide areas in the counting. Given  the very large number of craters comprised within a single region restricted to the northern part of a HiRISE image, next studies will involve machine learning algorithms (e.g. [10]) to perform automatic segmentation and mapping of the impact craters. These tools will be mandatory to refine our results over larger regions, and reconcile local and global scale analyses of the Martian surface.

Figure 1: Log-log plot of depth d versus diameter D for the crater dataset on the landslide. Plain black points are included in the linear regression (red dashed line). The gray points are other craters from the dataset. The scaling law is displayed under the logarithmic form, with r² value to quantify the regression robustness.

References: [1] McEwen A. S. et al. (2007) JGR: Planets, 112, E5. [2] Kirk R. L. (2008) JGR: Planets, 113, E3. [3] Boyce J. M. and Garbeil H. (2007) Geophysical Research Letters, 34, 16. [4] Garvin J. B. et al. (2003) Sixth International Conference on Mars, Abstract #3277. [5] Watters W. A. et al. (2015) JGR: Planets, 120, 2, 226-254. [6] Kite E. S. and David P. M. (2017) Icarus, 286, 212-222. [7] Quantin C. et al., (2004) Icarus, 172, 2, 555-572. [8] Millot C. et al. (2003) Submitted to Tenth International Conference on Mars, Abstract #3355 [9] Mangold N. et al. (2008) Icarus, 194, 2, 519-543. [10] C. Lee (2019), PSS, 170, 16-28.

How to cite: Millot, C., Quantin-Nataf, C., Dehouck, E., Volat, M., and Torres, I.: Morphometric Laws for Small Martian Craters (D < 50 m): a Case Study on Landslide Deposits, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-844, https://doi.org/10.5194/epsc2024-844, 2024.

How to cite: Hoyo García, M., Bruzzone, L., and Bovolo, F.: Enabling large-scale detection of subsurface features on Mars by their radar mapping, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-61, https://doi.org/10.5194/epsc2024-61, 2024.

Introduction

China’s first Mars exploration mission, Tianwen-1, was successfully launched on 23rd July 2020. The mission included an orbiter, a lander, and the Zhurong rover. Zhurong landed on Mars in the southern part of Utopia Planitia (109.925°E, 25.066°N) on 15th May 2021 and then conducted unprecedented in-situ observations with several instruments, including a ground-penetrating radar (GPR) called Rover Subsurface Penetrating Radar (RoSPR) (Liu et al., 2023). The radar (Fig. 1) aimed at probing the subsurface at the Zhurong landing site and searching for buried water ice. Utopia Planitia is well known for showing different ice-related landforms (Séjourné et al., 2019). To achieve this goal, RoSPR was equipped with two channels: a high-frequency channel providing high-resolution (5 cm, when permittivity is 3) full-polarization radargram of the shallow subsurface (first 10 meters) and a low-frequency channel capable of penetrating the Martian subsurface down to a depth of ~100 meters with a 1-meter vertical resolution.

In-situ GPR observation can reveal the structure and permittivity distribution along the cross-section drawn by the rover’s traverse. The subsurface structure is interpreted from the radar profile, which intuitively shows the location of subsurface reflectors, such as buried craters, aeolian deposits, and icy beds. Permittivity is a complex parameter describing the electromagnetic properties of materials. It mainly depends on composition and porosity. Its real part and imaginary part respectively control the wave velocity and signal attenuation rate in the Martian subsurface.

Data and Methods

We use all the 5893 dual-channel radar traces collected by RoSPR during the 325 solar days on Mars. The traverse of the radar is shown in Fig. 2. Those observations covered ~1840 meters with an elevation variation of ~10 meters from the Zhurong landing site southward. Data from ground experiments of RoSPR on glaciers on Earth are also used to provide a comparison between the RoSPR data acquired on different sites.

In addition to conventional RoSPR data treatments, advanced data processing methods are deployed to better reveal subsurface structure. In particular, we adopt a background matrix method (Rashed and Harbi, 2014) which statistically removes the systematic clutters. Blind deconvolution methods of the radar B-scan are also tested to address and eliminate the ringing effect that commonly exists. In addition, the bandwidth extrapolation technique designed for range resolution enhancement (Oudart et al., 2021) will be applied to further provide clarity among different buried objects.

Preliminary Results and Ongoing Work

To estimate the Martian soil permittivity, we applied a hyperbola fitting method which accounts for the antenna’s height (Wang et al. 2021, Oudart et al., 2022), and a centroid frequency shift method (Quan et Harris, 1997) which uses the real-time spectrum of dual-channel RoSPR data to maximize the accuracy (Fig. 3). Loss tangent estimation also reveals upper bounds of ~0.008 and ~0.022 for the first 10 meters and 40 meters, respectively.

The low-frequency B-scan, as shown in Fig. 4, reveals buried structures and indistinct reflectors in deeper regions. The variation of the B-scan and loss tangent value along the depth implies varied composition at the landing site. The ongoing work focuses on refining data processing techniques to improve the clarity of subsurface imaging and provide robust interpretation and geological analysis.
 References

[1] S. Liu et al., “Data Pre-Processing and Signal Analysis of Tianwen-1 Rover Penetrating Radar,” Remote Sensing, vol. 15, no. 4, p. 966, Feb. 2023, doi: 10.3390/rs15040966.

[2] N. Oudart et al., “Range resolution enhancement of WISDOM/ExoMars radar soundings by the Bandwidth Extrapolation technique: Validation and application to field campaign measurements,” Planetary and Space Science, vol. 197, p. 105173, Mar. 2021, doi: 10.1016/j.pss.2021.105173.

[3] N. Oudart et al., “Retrieval of the ground dielectric permittivity by planetary GPR accommodated on a rover: Application to the estimation of the reflectors’ depth by the WISDOM/ExoMars radar,” Planetary and Space Science, vol. 224, p. 105606, Dec. 2022, doi: 10.1016/j.pss.2022.105606.

[4] Y. Quan and J. M. Harris, “Seismic attenuation tomography using the frequency shift method,” GEOPHYSICS, vol. 62, no. 3, pp. 895–905, May 1997, doi: 10.1190/1.1444197.

[5] M. Rashed and H. Harbi, “Background matrix subtraction (BMS): A novel background removal algorithm for GPR data,” Journal of Applied Geophysics, vol. 106, pp. 154–163, Jul. 2014, doi: 10.1016/j.jappgeo.2014.04.022.

[6]A. Séjourné et al., “Grid Mapping the Northern Plains of Mars: Using Morphotype and Distribution of Ice-Related Landforms to Understand Multiple Ice-Rich Deposits in Utopia Planitia,” Journal of Geophysical Research: Planets, vol. 124, no. 2, pp. 483–503, Feb. 2019, doi: 10.1029/2018JE005665.

[7] R. Wang et al., “A Novel Approach for Permittivity Estimation of Lunar Regolith Using the Lunar Penetrating Radar Onboard Chang’E-4 Rover,” Remote Sensing, vol. 13, no. 18, p. 3679, Sep. 2021, doi: 10.3390/rs13183679.

Fig. 1 The dual-channel RoSPR onboard Zhurong rover. The low-frequency monopole antennas are mounted in the front of the rover. The high-frequency Vivaldi antennas are tied to the rover’s body.

Fig. 2 The traverse of Zhurong rover and the elevation at the landing site.

Fig. 3 The loss tangent estimation via RoSPR low-frequency data. The figure demonstrates the centroid frequency decay versus time delay relationship, calculated from 160 traces of Low-frequency channel data ranging from ~300 to ~550 ns.

Fig. 4 RoSPR low-frequency B-scan section from sol 112 to sol 288.

How to cite: Zhang, Z., Schmidt, F., Su, Y., Oudart, N., Le Gall, A., Saintenoy, A., Brighi, E., Ciarletti, V., Léger, E., Costard, F., and Séjourné, A.: Analysis of the Martian subsurface from the observation of the Rover Subsurface Penetrating Radar (RoSPR) onboard Tianwen-1, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-441, https://doi.org/10.5194/epsc2024-441, 2024.

Figure 2 Radargrams give the reflected signal where white is high reflecting amplitude and black low amplitude. The depth axis for this radargram is 35 meters.

Geological Modeling: The RIMFAX data can be used to make 3D models of the subsurface. The depth and locations of selected layers are picked from the RIMFAX Marginal Unit radargrams. From these layers, a geological model is made using GemPy, see [4] which employs a universal cokriging interpolation method, [4]. GemPy allows the user to generate complex 3D structural geological models through the interpolation of layer interfaces and orientation measurements and topography.  

Figure 3 shows a 2D cross section through the GemPy model. Figure 4 shows a map view of the model, indicating locations of outcropping layer boundaries. Figure 5 shows a 3D view of the subsurface layers in relation to surface topography.

Figure 3. 2D cross section through the GemPy model from West to East. Non-shaded layers are truncated by surface topography.

Figure 4. GemPy model reconstructed geologic map based on the extrapolation of picked layer depths and locations in the RIMFAX radargrams in the Margin Unit (dots).   

Figure 5. 3D geological model reconstructed from RIMFAX data showing the subsurface dipping layers and where they crop out on the surface.

Discussion: The large-scale subsurface structures revealed by the RIMFAX in the Marginal Unit are reminiscent of deltaic clinoforms. Figure 4 shows that Marginal Unit layers are cropping out in complicated patterns that are not apparent in orbital images or topographic maps. Deeper eroded layers are predicted to be cropping out in the wall of Neretva Vallis. Using RIMFAX radar data in conjunction with geological modelling is a powerful tool for studying subsurface structures on Mars.

Acknowledgments: The data used in this work are available at the NASA PDS Geosciences Node (https://doi.org/10.17189/1522644). This work was supported by the Research Council of Norway, grant no. 309835.

References:

[1] Farley et. al. (2021) Space Science Review, [2] Hamran et. al. (2021) Space Science Review, [3] Hamran et. al. (2022) Science Advances, [4] de la Varga et. al. (2019) GemPy, Geosci. Model Dev., 12, 1–32

How to cite: Hamran, S.-E., Paige, D., Andersson, F., Berger, T., Cardarelli, E., Carter, L., Dypvik, H., Russell, P., Mellon, M., Nunes, D., and Plettemeier, D.: Geological Modeling of Subsurface Structures from RIMFAX Ground Penetrating Radar Observations in Jezero Crater, Mars.  , Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-403, https://doi.org/10.5194/epsc2024-403, 2024.

Introduction :

In planetary sciences, accessing the subsurface is essential to fully understand the geological history of a given site. GPR (Ground Penetrating Radar) offers a way to investigate the subsurface in a non-destructive manner. So far, five GPR have investigated the subsurface of the Moon and Mars from the surface. RIMFAX GPR [1] onboard the Perseverance rover of NASA’s Mars 2020 mission is still operating at Jezero crater and RoSPR [2] investigated the subsurface of Utopia Planitia along the traverse of Zhurong, the rover of the CNSA’s Tianwen-1 mission. Both RIMFAX and RoSPR are ultra-wide band GPR ; RIMFAX is a Frequency Modulated Continuous Wave radar operating from 150 to 1200 MHz while RoSPR is equipped with two channels : which respectively operate in the 15-95 MHz (low frequency or LF channel) and the 450-2150 MHz (high frequency or HF channel) ranges. The next GPR to Mars should be WISDOM [5] onboard the Rosalind Franklin rover of the ESA’s ExoMars mission [8]; it is designed to probe the subsurface at Oxia Planum at frequencies in the 0.5-3 GHz range.

Clear buried interfaces or resolvable large reflecting structures are probably the most easily interpretable features a GPR can detect in the subsurface. However, radargrams often display diffuse scattering in the subsurface volume rather than clear structures. As an example, the radargram acquired by RoSPR at Utopia Planitia on Mars points to the presence in the subsurface of two layers containing scatterers of different size or nature, but no clear interface [3]. Scattering signature was also observed in the lunar subsurface by the LPR (Lunar Penetrating Radar), and the typical size of heterogeneities in the lunar regolith has been constraint [4].

In this work we propose to apply on RoSPR and, to a lesser extent, RIMFAX data a method developed on WISDOM simulated data [5]. This method aims at statistically retrieving the typical size of buried heterogeneities or scatterers from scattered signals in radargrams. The RoSPR LF profile section selected as an example to illustrate the method is shown on Fig 1. The orange section is 280 m long and corresponds to the end of the profile published in [3] displaying unambiguous scattering down to depths of 50 meters.

Method :

We conducted FDTD simulations with the TEMSI-FD code [6], that accurately models the WISDOM antennas in order to reproduce reliable soundings over 3-D synthetic heterogeneous media as illustrated by Fig 2. We generated a statistical number of subsurface volumes to perform meaningful statistical averages. Fig. 2 upper panel shows three examples of modelled subsurfaces with different typical size of scatterers L. Volume scattering occurs for each of the three L values considered, but the strongest volume scattering is not observed at the same frequencies. In other words, the wide frequency bandwidth of WISDOM allows to study the scattering process by sub-frequency bands to estimate the typical size of the buried structures. Fig. 3 synthesizes the results obtained with a sliding frequency window analysis for L values in the range 0.7-6.2 cm. A previous study [7] performed on simulated data shows that the maximum volume scattering intensity occurs for a wavelength λ = (5.3±0.2)L regardless of the value of L. We propose to use this result on Martian radargrams displaying diffuse scattering to constraint the L value in the investigated areas.

We built spectrograms applying a Fourier Transform to have a quick and global picture of the frequency content along the profile. This can be done on selected time windows of the radargram that correspond to given depth ranges (see Fig 4). This method also reveals possible horizontal variability in the frequency content of the radar signal. The described analysis requires in the first place to compensate the variability of the instrument’s gain over the bandwidth . This is achieved by dividing the frequency domain data by the average RoSPR spectra of the subsurface of Utopia Planitia (a processing called whitening). However, in absence of proper calibration data, this correction is not perfect and some instrumental artifacts may remain (as it is the case on Fig. 4 around 47 and 84 MHz).

Results and preliminary analysis :

The result shown in Fig. 3 implies that ultra-wide band GPR data can be used to determine the size of heterogeneities within a range that depends on its frequency range of operation and of the average permittivity of the medium, as synthesized in Fig 5 for all Martian GPRs.

In Fig. 6 the spectrograms’ frequency-axes corresponding to different time windows (Fig 4) are converted into L-axes via the relation established based on simulations (Fig. 3) in order to interpret them in terms of scatterer size. The main features are indicated in Fig 6, but must be taken with cautious  In particular, peaks in the volume scattering intensity can have various origins and only global maximums should be considered. Another caveat is the possible presence of several populations of scatterers of different sizes.

The difference between upper and lower maps in Fig. 6 somewhat supports the interpretation advanced by [3] with two layers characterized by smaller scatterers in the upper layer than in the lower layer). The spectrogram analysis method could put firmer constraints on their actual size distributions and horizontal scales of variability.

References :

[1] Hamran et al., Space Sci Rev (2020). https://doi.org/10.1007/s11214-020-00740-4

[2] Liu et al., Remote Sensing 2023. https://doi.org/10.3390/rs15040966

[3] Li et al., Nature (2022). https://doi.org/10.1038/s41586-022-05147-5

[4] Ding et al., A&A 2022. DOI: 10.1051/0004-6361/202142803

[5] Ciarletti et al., Astrobiology 2017. https://doi.org/10.1089/ast.2016.1532

[6] Ciarletti et al., Monthly Notices of the Royal Astronomical Society, 2017. https://doi.org/10.1093/mnras/stx3132

[7] Brighi et al., Planetary and Space Science, under review

[8] Vago et al., Astrobiology, 2017. https://doi.org/10.1089/ast.2016.1533

How to cite: Brighi, É., Oudart, N., Ciarletti, V., Le Gall, A., Zhang, Z., Su, Y., Schmidt, F., Saintenoy, A., and Hamran, S.-E.: Spectral analysis of radargrams of the Martian subsurface to constrain scatterer sizes-Application to RoSPR/Zhurong and RIMFAX/Perseverance GPR data, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-439, https://doi.org/10.5194/epsc2024-439, 2024.

The ever-changing transparency of the martian atmosphere hinders the determination of absolute surface colour from spacecraft images. While individual high-resolution images from low orbit reveal numerous colour details of the geology, the colour variation between images caused by scattering off atmospheric dust can easily be of greater magnitude. The construction of contiguous large-scale mosaics has thus required a strategy to suppress the influence of scattering, often a form of high-pass filtering, which limits their ability to convey colour variation information over distances greater than the dimensions of single images. Here we use a dedicated high altitude observation campaign with the Mars Express High Resolution Stereo Camera (HRSC), applying a novel iterative method to construct a globally self-consistent colour model. We apply the model to colour-reference a high-altitude mosaic incorporating long-range colour variation information. Using only the relative colour information internal to individual images, the influence of absolute image to image colour changes caused by scattering is minimised, while the model enables colour variations across image boundaries to be self-consistently reconstructed. The resulting mosaic shows a level of colour detail comparable to single images, while maintaining continuity of colour features over much greater distances, thereby increasing the utility of HRSC colour images in the tracing and analysis of martian surface structures.

How to cite: Michael, G., Tirsch, D., Matz, K.-D., Zuschneid, W., Hauber, E., Gwinner, K., Walter, S., Jaumann, R., Roatsch, T., Postberg, F., and Liu, J.: A global colour mosaic of Mars from Mars Express HRSC high altitude observations, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1166, https://doi.org/10.5194/epsc2024-1166, 2024.

Introduction :

Martian H₂O, CO₂ and dust cycles are key processes in modern climate dynamics. The polar caps [1], as well as their icy outliers (that can take the form of convex-shaped ice mounds inside deep craters) are the current, major, exposed volatiles reservoirs. Unanswered questions remain about the mass balance of the caps. Concerning outliers, such as the ice domes of Louth and Korolev craters[2,3], the prevailing hypothesis suggests a build up by atmospheric deposition [4]. If accurate, these icy mounds could record past climate  conditions while still responding to present-day Martian climate change because of their southerly location [5].  They therefore offer unique opportunities to study volatiles dynamics, with different boundary conditions than those of the north or south caps.

Research question :

Our study takes advantage of the Exomars Trace Gas Orbiter – Colour and Stereo Surface Imaging System  (CaSSIS) [6] colour imaging capabilities to evaluate pluri-annual ice changes, especially water frost and dust coverage at Louth and Korolev craters ice mounds. Those changes can serve as a proxy of local dynamics, which can reveal important information about ice and dust cycles.

General methods:

CaSSIS is the high-resolution stereoscopic camera of ESA’s ExoMars Trace Gas Orbiter (TGO). It possesses 4 broad band filters across the visible range : BLU (497nm), PAN (677nm), RED (835nm) and NIR (940nm). The orbit of TGO being non-Sun-synchronous, it allows us to acquire images of the same region of interest at different seasons and local times, at 4.6 m/pixel with high a  signal-to-noise ratio, and therefore allows to monitor diurnal phenomena until 74°N. The first step of the study consists of a morphological and spectral analysis of 25 Louth crater images, that cover MY34 to MY37.

Visible reflectance features of water ice include an average high reflectance, relatively constant over the VIS range, as well as asymmetric overtone absorption bands at 800nm, 840nm, and 1030nm. In CaSSIS colors, the strongest signature of water ice is the reduction of the strong visible red slope of the dust. It has been demonstrated experimentally that CaSSIS color criteria can help differentiate  between different types of ice and dust mixings [7].

Observational results  : 

In accordance with previous studies, [5,8], Louth ice mound appears generally homogeneous, with little to no significant spatial variations in its global albedo and colour. This is not the case in two particular observations, taken on 15.02.2022, and on 28.03.2022, respectively. A dark pattern appeared and displayed changes in the eastern direction (in Fig.1, the red circle shows the area of change). 

Fig.1 : Two CaSSIS images (MY36_018836_078 and MY36_018951_106) of Louth crater, 

taken in early martian autumn. The two images are NIR/PAN/BLU composites. 

We searched for reports of similar events in the literature, and only found one study [9] reporting an ‘unexplained dark patterning evolution’ in MRO - CTX images. We found other occurrences of dark patterns in MRO - HiRISE [10], in 2010, 2022, and 2023. Fig. 2 displays two examples of observations, both with HiRISE and CTX, in 2022 and 2023.

Fig.2: Two occurrences of dark patterns in 2022 (A) and 2023 (B)

The CTX observation shown in Fig.2 A was taken 14 days before the CaSSIS image MY36_018836_078 shown in Fig 1. A notable  eastward progression can be distinguished.Each occurrence is different in shape, shows  strong seasonality (appears at the end of summer, evolves through the autumn equinox), as well as rapid evolution (~tens of Martian days). The dark patterns disappear during each wintertime, showing a white  mound each following spring. This indicates that the mechanism that produces these patterns is related to seasonal effects.

Preliminary spectral measurements of the ice in CaSSIS images show that the signal is largely dominated by the dust signature. We hypothesize that the dark marks could be caused by either deposition from the adjacent dune field, or local sublimation events.

What’s next ?  

While continuing on spectral analysis with CaSSIS, we would like to explore the NOMAD [11] dataset to search if outliers in atmospheric water vapor content are correlated with the dark patterns occurrences. A ‘cloud’ can be seen in the MY36_018951_106 observation displayed in Fig.1. We do not know if this is linked to the formation of the North Polar Hood, or to the event itself. Thirdly, we will make use of the Mars Climate Database [12] to explore the temperature and relative humidity fields around Louth. Lastly, we will perform a similar study on Korolev crater. It is less clear than in Louth, but some structures appearing in early autumn resemble the ones that have been described here.

Acknowledgments:

This work has been developed in the framework of the National Center for Competence in Research PlanetS funded by the Swiss National Science Foundation (SNSF).CaSSIS is a project of the University of Bern and funded through the Swiss Space Office via ESA's PRODEX programme. The instrument hardware development was also supported by the Italian Space Agency (ASI) (ASI-INAF agreement no. 2020-17-HH.0), INAF/Astronomical Observatory of Padova, and the Space Research Center (CBK) in Warsaw.

References:

[1]  Leighton,  et al.. Sci. Vol.153, (1966).

[2]  Brown, A J. et al. Icar., Vol.196, (2008). 

[3]  Armstrong, J.C, et al. Icar. Vol. 174 (2005).

[4]  Conway, S. et al. Icar. Vol .220. (2012).

[5]  Bapst, J. et al. Icar., Vol. 308,(2018)

[6]  Thomas N et al. Space Sci. Rev. Vol. 212 (2017)

[7]  Yoldi,Z. et al,Icar,Vol.358, (2021)

[8]  Bapst J , Byrne S  LPSC, Abstract #6097  (2015)

[9]  Brown AJ et al. arXiv:1711.06372 (2017)

[10]  McEwen,A.S., et al. Journal of Geo. Research: Planets 112.E5 (2007).

[11]  Vandaele, AC, et al. Pl. and Space Sci. 119 (2015)

[12] Millour, E, et al. , EPSC2015-438, Vol. 10 (2015)

How to cite: Azevedo (1), M., Pommerol (1), A., and Thomas (1), N.: Exploring the seasonal variability at Louth crater ice mound with orbital colour imaging , Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-611, https://doi.org/10.5194/epsc2024-611, 2024.

The 2-decade long observation time series of the High Resolution Stereo Camera (HRSC) [1] onboard Mars Express provides the precious opportunity to search for long-term surface changes on Mars in comparison to other multi-temporal image data of even earlier missions. Previous works on the topic of Martian change detection have mostly focused on either different sensor types (e.g., CaSSIS data [2] and MCC data [3]) or on short term and small-scale processes (e.g., gully formation [4] and impact cratering [5]). Surface changes on Mars that are not related to seasonal processes or are often too small-scale to be observed with medium- or low-resolution images occur very slowly. Besides impact cratering, aeolian erosion and deposition are the most important geological processes that cause perceptible surface changes today. As the winds on Mars have a low transport capacity due to the thin atmosphere, slow large-scale changes can best be observed in images with long time intervals. The identification of such long-term surface changes will help gaining a better understanding of Mars’ surface-atmosphere interactions and its general surface evolution.

This work starts out by visually comparing Viking orbiter images (from 1975 to 1980) and HRSC images (from 2004 to today) to find changes. We will additionally use image data from other sensors, such as CTX and MOC, to include observation with different spatial resolutions and to narrow down the time of the possible change process. The changes to be expected are wind streak, sand sheet and dust deposition (albedo) changes as well as long-term changes of the polar caps, slope-streak development and larger dune migration. As a start, this work focusses on wind streak changes, which are easy to detect and confirm. We found a significant amount of modified wind streaks around impact craters with changing orientation (azimuth) and/or changes in size. We will present first results from dark and bright (Figure 01) wind streaks in various regions on Mars, which indicate both depositional and erosional wind activity. Additionally, we will use the iron oxide ratio, a remote sensing index that uses the ratio of the red and blue wavelengths to visualize oxidized material on the surface [6], to visualize changes in depositional activity between HRSC images from different acquisition times.  In a later step, we try to link the observed changes to specific atmospheric events that caused them (e.g., global dust storms). An overarching objective of our work is to assess the usability of HRSC data for detecting changes on the Martian surface and to provide a dataset of multi-temporal images that can be used for future studies and analysis of specific areas or features of interest.

Figure 01: Comparison of bright wind streaks in Syrtis Major, where we already found a significant amount of wind streaks with changing azimuth and/or changing size. Image A shows the area in the Viking MDIM2.1 Grayscale Global Mosaic with a resolution of 232 m. Image B shows the area in the HRSC image hm612_0000 with a resolution of 200 m. Sub-image a and b show the corresponding detailed section for a better analysis of the changing azimuth and/or size of the three wind streaks.

References

[1] Jaumann, R., et al., (2007). The high-resolution stereo camera (HRSC) experiment on Mars Express: Instrument aspects and experiment conduct from interplanetary cruise through the nominal mission. Planetary and Space Science 55, 928-952.

[2] Rangarajan, V., Tornabene, L., Osinski, G., Conway, S., Seelos, F., Silvestro, S., . . . Cremonese, G. (2023, January 26). Change detection and monitoring of active Martian surface phenomena. Icarus.

[3] Misra, I., Rohil, M., Moorthi, S., & Dhar, D. (2021). Mars Surface Multi-decadal Change Detection Using IRSO's Mars Color Camera (MCC) and Viking Orbiter Images. Springer Nature Singapore Pte Ltd.

[4] Pasquon, K., Gargani, J., Massé, M., & Conway, S. (2016, April 4). Present-day formation and seasonal evolution of linear dune gullies on Mars. Icarus.

[5] May, S. (2018). METEOR IMPACT DETECTION ON MARS WITH CHANGE DETECTION FRAMEWORK.

[6] Segal, D. "Theoretical Basis for Differentiation of Ferric-Iron Bearing Minerals, Using Landsat MSS Data." Proceedings of Symposium for Remote Sensing of Environment, 2nd Thematic Conference on Remote Sensing for Exploratory Geology, TX (1982): pp. 949-951." (ESRI, 2018)

How to cite: Schöpf, A., Tirsch, D., Jaumann, R., and Matz, K.-D.: Long Term Change Detection on Mars using HRSC images in comparison to other multi-temporal image data sets, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-201, https://doi.org/10.5194/epsc2024-201, 2024.

We are deriving insights about the climate evolution of Mars by automating surface mapping and monitoring with AI models. We are also analysing these models’ uncertainty and decision-making process with explainable AI and uncertainty estimation techniques. We focus on ice block fall detection at the north polar region and polygon detection globally.

The north polar ice cap on Mars is comprised of ice layers relating to the climate cycles of Mars and this way it preserves a ∼ 4 million years climate change record of the planet. These layers are exposed at the marginal steep scarps of the cap, that are currently active with avalanches and block falls. We are monitoring the mass wasting activity at the north polar region by detecting ice blocks [1] and their sources [2]. We estimate the current erosion rates of all scarps resulting in a detailed map of how this extensive ice-layered dome is being shaped today.

At the mid-latitude regions, where large volumes of excess ice exist, young patterned ground resembles glacial and periglacial patterns on Earth. Can freeze-thaw cycles have recently thawed the permafrost on Mars to produce these landforms? This would have implications for the recent hydrologic past of the planet. We detect young ice-wedge polygons to determine their distribution and relationship to the topography with the potential of elucidating the formation mechanism and the role of liquid water in the recent past of Mars.

We use AI models to automate surface mapping and monitoring, because they outperform all other methods. However, they are treated as black box systems. We want to know why a model produces a specific response and how certain it is about the correctness of each result. To answer these questions, we put together an application-independent framework that deploys uncertainty estimation and explainable AI methods to provide insights into the decision-making process and assess the uncertainty of the results [3], in this project the uncertainty of the surface mapping.

References:

[1] Martynchuk O. et al., Computer vision model for monitoring mas wasting activity in the Martian North Polar region. EPSC 2024.

[2] Su S. et al., Current mass wasting on icy scarps of Mars: A comprehensive mapping of ice-fragments in the north polar area. AGU 2023.

[3] Fanara L. et al., W²: How Certain and Why? Uncertainty Estimation and Explainable AI Applied on Mars Projects. AGU 2023.

How to cite: Fanara, L., Su, S., Martynchuk, O., Hauber, E., and Gwinner, K.:  AI-driven surface mapping and explainable AI for climate insights on Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-948, https://doi.org/10.5194/epsc2024-948, 2024.

Previous research has suggested that large floods in the Elysium mons region, Mars, may have melted significant amounts of ice in the late Amazonian epoch.  Recent work has proposed that there may be remnants of ice-rich layers in the subsurface of the Medusa Fossae Formation (Watters et al., 2024), close to Elysium mons, see Figure 1.  In this work we investigate the plausible presence of ices in Athabasca Valles region, which is located at the southern flank of Elysium Mons and close to the Medusa Fossae Formation (MFF). 

Figure 1. Situation map of the area of interest considered in this study.

Athabasca Valles consists of a group of subparallel valleys with NE-SW orientation and SW vergence whose source area is located at 10ºN, 157ºE. This area has a consistent slope that stops at Cerberus Palus, see Figure 1. There are currently two different models to explain the recent geological history of Athabasca Valles: one of them postulates the existence of considerable amounts of ice in the subsurface that interacted with magma (Cassanelli and Head, 2018), whereas  the other one considers only volcanic activity (Ryan and Christensen, 2012; Miller et al., 2023).

We have investigated the Athabasca Valles region to evaluate the possible presence of ice in the recent past. Using MOLA, THEMIS, HRSC, CTX and HiRISE observations, we find grooved surfaces at the bottom of the valleys, hills with nunatak like shapes, potential drop stones, surface features like polygonal patterned floors and pingo-like landforms, see Figure 2, as well as some other periglacial related landforms (Balme and Gallagher, 2009; Yakovlev et al., 2021). All of these features are caused by the sustained presence of ice.

Figure 2. CTX and HiRISE view. Top left: Eroded impact crater surrounded by grooved surfaces and possible dropstones, in the MurrayLab_CTX_V01_E152_N08 tile. Top right: Pingo like features and polygonal patterned floors in the HiRISE HI_046690_1900_063437_1900 image. Bottom left: Possible moraines and grooved surfaces observed in MurrayLab_CTX_V01_E156_N08 tile. Bottom right: nunatak-type mount observed in MurrayLab_CTX_V01_E156_N08 tile. All images have the same orientation.

We also observe the existence of a platy region with fractures that may be attributed to lava-ice interaction, see Figure 2. Using crater counting chronology methods, we have dated the platy surface structures and estimated an age of 29 Ma.  We propose that the existence of ices under the surface 29 Ma ago caused the movement of the platy surface structures downstream of Athabasca Valles because of the lava-ice interaction. The valley has a consistent slope that stops at Cerberus Palus, and it is reasonable to think that ice or melted ice may have moved along this slope accumulating at Cerberus Palus.

Figure 3. THEMIS Thermal image of a fractured plate structure, where the plates show low thermal inertia whereas the inter-plate regions show high thermal inertia. These two differentiated parts have been dated using crater counting methods (see inset).

Using SHARAD data, we find bright radar reflectors in the accumulation area of Cerberus Palus, which we interpret as potential preserved subsurface layers of ice pore-filling materials, remains of previously existing ices, see Figure 4. We have modelled the temperatures of an hypothetical ice sheet under the surface of this region using the same methodology described by Egea et al. 2022 (Egea-González et al., 2022) for the North Pole region.  Our analysis shows that subsurface ice can be thermally stable under the present-day thermal environment.

Our observations are consistent with the possible interaction of lava and ices in Athabasca Valles around 29 Ma ago, and the present-day preservation of ice in the subsurface at the end of the valley, in Cerberus Palus.

Figure 4. Example of one radargram where we can observe the existence of subsurface reflectors at the end of the basen.

Acknowledgements

M.-P.Z. and F. M. were supported by grant PID2022-140180OB-C21 funded by MCIN /AEI /10.13039/501100011033 / FEDER, UE. The work of F. Mansilla is supported by the grant PRE2020-09170, founded by MCIN/AEI/10.13039/501100011033.

References

Balme, M. R. and Gallagher, C.: An equatorial periglacial landscape on Mars, Earth Planet. Sci. Lett., 285, 1–15, https://doi.org/10.1016/j.epsl.2009.05.031, 2009.

Cassanelli, J. P. and Head, J. W.: Large-scale lava-ice interactions on Mars: Investigating its role during Late Amazonian Central Elysium Planitia volcanism and the formation of Athabasca Valles, Planet. Space Sci., 158, 96–109, https://doi.org/10.1016/j.pss.2018.04.024, 2018.

Egea-González, I., Lois, P. C., Jiménez-Díaz, A., Bramson, A. M., Sori, M. M., Tendero-Ventanas, J.-Á., and Ruiz, J.: The stability of a liquid-water body below the south polar cap of Mars, Icarus, 383, 115073, https://doi.org/10.1016/j.icarus.2022.115073, 2022.

Miller, R. C., Grima, C., Gulick, S. P. S., Goudge, T. A., Russell, A. T., Perry, M. R., Putzig, N. E., and Campbell, B. A.: Dynamic development of the Athabasca Valles outflow system from volcanic facies and 15 m scale roughness, Icarus, 115691, https://doi.org/10.1016/j.icarus.2023.115691, 2023.

Ryan, A. J. and Christensen, P. R.: Coils and Polygonal Crust in the Athabasca Valles Region, Mars, as Evidence for a Volcanic History, Science, 336, 449–452, https://doi.org/10.1126/science.1219437, 2012.

Watters, T. R., Campbell, B. A., Leuschen, C. J., Morgan, G. A., Cicchetti, A., Orosei, R., and Plaut, J. J.: Evidence of Ice‐Rich Layered Deposits in the Medusae Fossae Formation of Mars, Geophys. Res. Lett., 51, e2023GL105490, https://doi.org/10.1029/2023GL105490, 2024.

Yakovlev, V., Horelik, S., and Lytvynenko, Y.: On the cryogenic nature of the large hills of Mars, Planet. Space Sci., 208, 105340, https://doi.org/10.1016/j.pss.2021.105340, 2021.

How to cite: Mansilla Nuñez, F., Zorzano, M.-P., and Ruiz, J.: Potential existence and preservation of subsurface ice in the Athabasca Valles region, Mars., Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-361, https://doi.org/10.5194/epsc2024-361, 2024.

Introduction 

Orbital and in situ images have revealed past surface water activity on Mars. Morphological observations include valley networks [1, 2] and the presence of a few hundred paleolakes [3-5]. Jezero crater, which is the landing site for the Perseverance rover, is one of the most extensively studied palaeolakes today. This site is of interest for understanding ancient martian hydrologic activity, and for finding potential traces of past life. Current observations suggest that Jezero crater was fed by two inlet valleys [6] which are at the origin of two sedimentary fan deposits, interpreted as deltaic deposits [6, 7]. An outlet valley, Pliva Vallis, is present on the eastern crater rim [6] raising the question of whether the lake system operated mainly as an open basin, or as a closed basin system with episodic overflows. Resolving this uncertainty would improve our understanding of the Pliva Vallis’ evolution within time and consequently our knowledge of the history of Jezero’s lake.

Observations and interpretation

We performed a qualitative morphological study of the outlet valley using digital elevation models and orbital images (CTX and HiRISE data) acquired by Mars Reconnaissance Orbiter. First, we observed stratified and re-incised sedimentary deposits on the floor and inner walls of Pliva Vallis that we infer are a result of the fluvial activity of this valley. Alternating fine deposits and strata containing meter-scale boulders reflect the variation of flow intensity over time. Regarding the morphometry of the outlet valley, we observed several topographical features that are atypical for overland flows [8].

Indeed, the decrease of the valley width and depth from upstream to downstream, the presence of a perched channel and the observations of several topographic rises along the valley floor suggest a valley with a discontinuous activity and at least two episodes of energetic overflow. In addition, a few meters from the start of Pliva Vallis, we noticed the presence of four mains topographical steps within the wall of the valley, which are marked by the four colored lines in Figure 1 and are interpreted as four bedrock incisions terraces. This morphological observation confirms the episodic activity of Pliva Vallis with four rapid flow events caused by four phases of breach in the eastern rim of Jezero crater.

These observations suggest a scenario of lake overflow with at least four different events, responsible for the progressive and discontinuous incision of the outlet valley. Thus, we propose that the Jezero’s crater lake overflowed sporadically during the four breach events before emptying completely as a closed system though evaporation and infiltration, with limited fluvial input [9].

Modeling of the breach erosion processes

To give a first estimate of the duration of these four episodes of lake overflow and valley incision, we use a simple 0-D mathematical model, simulating valley formation by breach erosion and crater lake overspilling [10, 11]. This model provides a simple simulation of breach formation, particle transport and erosion rate in the outlet valley during time. Considering the imposed assumptions, (e.g., uniform valley geometry, no pre-existing valley, no downstream change in flow velocity, etc.), we applied the model to the Jezero’s crater lake in order to simulate the four overflow episodes previously identified, using the valley width and depth and the local slope measured for each episode where the four erosional terraces are identified. We assume a flow into the crater of about 100 m³/s, and we found that varying this value has little effect on our results given the order of magnitudes differences in discharge rate between this steady-state process and the overflows. Figure 2A shows the decrease of particle transport (D50 of 0.076 m [12]) over time for each event and for a range of slopes. As soon as the breach occurs, the flow loses energy progressively at first and then rapidly, until it can carry no more particles. The erosion rate can also be calculated during lake emptying for each episode (Figure 2B). Over time, erosion slows and tends to become zero as the flow decreases and loses energy. Modeling results suggest that each flooding and incision event in the outlet valley lasted less than a few weeks, or even a few days for some episodes. The time scales calculated with this model are consistent with our interpretation that flow events in the outlet valley were rapid and discontinuous.

Conclusion and outlook

The study of the outlet valley of Jezero crater helps us to understand better the activity of this lake. Our observations allow us to propose a new scenario of evolution of the Jezero’s crater lake - a closed basin system which was episodically and temporary opened during at least four identified overflow events.

References

[1] Pieri, D. C. (1980) Science, 90, 895–897.

[2] Carr, M. H. (1987) Nature, 326, 30–35.

[3] De Hon, R. A. (1992), Earth, Moon and Planets, 56, 95–122.

[4] Cabrol, N. A. and Grin, E. A. (1999), Icarus, 142, 160–172.

[5] Wharton, R. A., et al. (1995), Journal of Paleolimnology, 13, 267–283.

[6] Fassett, C. I. and Head, J. W. (2005), GRL, 32, L14201.

[7] Schon, S. C. et al. (2012), PSS, 67, 28–45.

[8] Schumm, S. A. and Ethridge, F. G. (1994), SEPM Society for Sedimentary and Geology, book. 

[9] Mangold, N. et al., (2024) Constraints on Jezero paleolake history from its fluvial input, Abstract for The Tenth International Conference on Mars.

[10] Warren, A. O. et al. (2021), EPSL, 554, 116671.

[11] Holo, S. J. and Kite. E. S. (2019), 50th LPSC, Abstract #2481.

[12] Mangold, N. et al. (2024), JGR Planets, 129, e2023JE008204. 

How to cite: Villette, J., Mangold, N., Conway, S., Le Deit, L., and Kite, E.: The evolution of Pliva Vallis, the outlet valley of the Jezero's crater lake on Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-236, https://doi.org/10.5194/epsc2024-236, 2024.

Introduction

Hebrus Valles (HV) and Hephaestus Fossae (HF) are well-preserved Early Amazonian outflow channel systems carved into bedrock in SE Utopia Planitia, Mars (17-25 °N, 118-129 °E, Fig. 1). They exhibit a diverse set of morphologies indicative of formation by one or more liquid water outflow events (Christiansen, 1987; http://ntrs.nasa.gov/search.jsp?R=19870019962). However, little is known about their history, including both the origin and fate of the water and resulting sediments. This represents a significant gap in our understanding of geologic processes involving liquid water occurring in the Amazonian Period.

Thanks to extensive coverage by recent datasets, it is now feasible to study the evolution of the HV-HF outflow channel systems with an integrated analysis of diverse and complementary data, including high-resolution visible and infrared imagery mosaics. We performed chronostratigraphic geologic mapping of the HV-HF region followed by a detailed reconstruction of the sequence of geological processes and events based on the identification of morphological and thermophysical facies. Then, we employed impact crater statistical analysis to determine geologic unit and outflow channel ages.

Geologic mapping

The oldest unit in the study area is the Nepenthes flow unit, a complex layered volcanic and sedimentary unit that was previously identified in the adjacent Nepenthes Planum region (Skinner and Tanaka, 2018; doi:10.3133/sim3389). On top of it lies the Utopia lower unit, which also corresponds to a sedimentary unit identified older geologic maps (Tanaka, 2005; https://pubs.usgs.gov/sim/2005/2888/; Skinner and Tanaka, 2018; doi:10.3133/sim3389). Moving up in the stratigraphic column, we found three additional units: the Utopia lowland unit, the Utopia lumpy unit, and the Utopia lobate unit. Based on morphological facies analyses and the observation of clear and distinct stratigraphic contacts, we determined that the Utopia lumpy unit and the Utopia lobate unit originated from subsequent and distinct tsunami-like events (e.g., Rodriguez et al., 2022; doi:10.1038/s41598-022-18082-2) transporting sedimentary materials across Utopia Planitia. Instead, the Utopia lowland unit and the Utopia lumpy unit are not delimited by clear stratigraphic boundaries, suggesting that they may be different facies of the same chronostratigraphic unit. Three younger units overlie the Utopia units. In the NE corner of the geologic map lies the Elysium volcanic unit, which consists of volcanic flows related to Elysium Mons volcanic activity. To the south, we identified two additional volcanic and tectonic units: the Elysium platy unit and the Elysium chaos unit.

Figure 1: Chronostratigraphic map of the HV-HF region.

Age measurements

            We conducted impact crater size-frequency distribution (CSFD) analyses on each chronostratigraphic unit by recording the locations and diameters of approximately 10000 impacts. We find formation ages for the Utopia units ranging from 2.9 Ga to 3.4 Ga (Fig. 2), placing the ages of units carved by channels between the Late Hesperian and Early Amazonian epochs. We also found evidence of resurfacing events occurring across three of the Utopia units at 2.2-2.3 Ga and a younger resurfacing event occurring at 1.7 Ga (Fig. 2). We interpret these events to be a combination of erosion of the Utopia Planitia plain-forming terrains and possible accumulation of volcanic sediments mantling the eastern half of the study area, perhaps associated with a peak in volcanic activity around 2-2.5 Ga (Platz and Michael, 2011; doi:10.1016/j.epsl.2011.10.001). Stratigraphy tells us that the Hephaestus Fossae and Hebrus Valles channel systems are necessarily younger than the surfaces they cut through, indicating that the channels must be younger than ~2.9 Ga. CSFDs within HV and in its distributary region yield ages of 1.8 Ga and 1.9 Ga, respectively (Fig. 3). Similarly, buffered crater counting of HV and HF yields ages of 2.1 Ga and 1.8 Ga, respectively (Fig. 3). These results confirm that the channels formed in the second half of the Early Amazonian about 1.5 Gy after the formation of the Utopia units, meaning that surficial aqueous activity in Utopia Planitia was occurring well into the Amazonian period.

Figure 2: CSFD analyses for the Utopia units with model formation ages (red) and resurfacing events (blue).

Figure 3: CSFD analyses for the HV and HF channels showing their model formation ages.

Origin and evolution of the outflow channels

The HV and HF channels originate from two deep pits within the Utopia lower unit (Fig. 1). We identified several elongated ridges protruding through the floor of these pits, which we interpret as remnants of magmatic intrusions that cracked the cryospheric seal of an ancient pressurized aquifer, as previously proposed for morphologically similar features in Cerberus Fossae (Head et al., 2003; doi:10.1029/2003GL017135). The morphology of the pits and the proximal sections of the channels are entirely compatible with morphologies observed in flume experiments where subsurface water discharge from a pressurized aquifer cause erosion and formation of terraces and deep source depressions within poorly consolidated sediments (Marra et al., 2014; doi:10.1002/2014JE004701). We reconstructed the initial phases of pit formation and channel incision in Fig. 4.

The pressurization of the aquifer can be explained by two mechanisms: heating of the aquifer by a magmatic intrusion that then cracked the cryospheric seal (Delaney, 1982; doi:10.1029/JB087iB09p07739) and/or the presence of an hydraulic gradient in an extensive cryospherically-sealed aquifer (e.g., Palumbo and Head, 2020; doi:10.1029/2020GL087230.). We find that both mechanisms could have contributed to the water discharge, and that the magnitude of their individual contributions depended on the permeability of the aquifer, because low permeability is required to maintain pressurization in the case of magmatic heating. We also note that the magmatic driver of pressurization and cryospheric cracking implies that the water release must have been modulated by magmatic activity, and thus we find it likely that multiple long-lived water discharge events occurred in this region until magmatic activity waned.

Figure 4: Step by step diagram of initial pit and channel formation. (left) A cryospheric seal confines a deep aquifer. Polygonal fractures are already present on the surface. (center) Magmatic intrusions crack the cryospheric seal leading to ponding of water, the formation of pits, and initial downslope flow on the surface. (right) Continuing water release deepens the source pits and carves channels in HV and HF, often following pre-existing fractures.

How to cite: Nerozzi, S., Spurling, R., and Holt, J.: The Fluvial and Volcanic History of Hebrus Valles and Hephaestus Fossae on Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-648, https://doi.org/10.5194/epsc2024-648, 2024.

Abstract

We present results of experiments performed inside a low-pressure chamber designed to investigate whether the volume of mud changes when exposed to a Martian atmospheric pressure. Depending on the mud viscosity, we observe a volumetric increase of up to 30% at the pressure of ∼6 mbar. The reason is that the low pressure causes instability of the water within the mud, leading to profuse bubble formation that increases the volume of the mixture. This mechanism bears resemblance to the volumetric changes associated with the degassing of terrestrial lava or mud volcano eruptions caused by a rapid ascend along the conduit and associated pressure drop.

Introduction

Subtle mounds have been discovered in the source areas of Martian kilometer-sized flows and on top of summit areas of some domes (Fig. 1; [1-2]). These features have been suggested to be related to subsurface sediment mobilization, opening questions regarding their formation mechanisms. Previous studies hypothesized that they mark the position of feeder vents through which mud was brought to the surface [1-2]. Two theories have been proposed: (a) ascent of more viscous mud during the late stage of eruption and (b) expansion of mud within the conduit due to the instability of water under Martian conditions.

Fig. 1. Examples of Martian mounds and knobs.

Mud volcanoes on Earth release mud with a wide range of viscosities [3]; likewise similar variations are also expected in the mud extruded on other planetary surfaces.  So far, experimental studies have been conducted only for low viscosity mud under low pressure conditions [4-5], and the behavior of more viscous muds under such conditions remains unclear. To overcome this knowledge gap, we completed analog experiments, to investigate the behavior of various viscous muds under martian pressure conditions.

Methods

We performed a set of experiments using the Mars Simulation Chamber at the Open University (UK), in which we inserted a two-part test bed that contained a reservoir which was able to accommodate 600 ml of mud and a ∼10 cm thick layer of frozen icy-sandy infill as the surrounding (Fig. 2). The temperature of this infill was around -20°C before experiment started. No active cooling of the icy-sandy infill was used, hence the infill slowly warmed up. All experiments were completed before the temperature of the icy-sandy infill reached the melting point of the water ice. The temperature of the mud when poured into the container was either 0.5-3°C (further refer as “cold”) or 13-22°C (“warm”).

Figure 2. Schematic illustration showing the experimental setup inside the Mars Simulation Chamber.

Three different viscosities were used (Fig. 3) and they are further refer as “low”, “medium” and “high.”  The mud mixtures were prepared by mixing deionized water with 1% w/w of dissolved magnesium sulfate salts (MgSO4) and clay content varying depending on the required viscosity. Magnesium sulfate salt was added to the water to achieve Earth's average river salinity, which enables the suspension of submillimeter clay particles within low viscosity mixtures. During the experiment, the pressure was gradually reduced from 1 bar to 5-7 mbar to achieve the martian surficial pressure. Two different speeds of pressure drops were utilized, namely the pressure was reduced within a timeframe of minutes (rapid) or more than an hour (slow).

Figure 3. Viscosity curves of “low”, “medium” and “high” viscosity aqueous bentonite.

In order to quantify volumetric changes of the mud samples, we used semi-manual and automatized image analyses using the PIV (Particle Image Velocimetry; [6]) and photogrammetry methods.

Results and discussion

Results revealed a significant volume increase during the experiments with slow depressurization, high mud viscosity and low initial mud temperature (Fig. 4). The volumetric change occurs due to the formation of water vapor bubbles, which are temporarily trapped within the mud. This phenomenon occurs since the bubble buoyancy is insufficient to overcome the drag force within the viscous material. Hence, these bubbles remain trapped in the mud allowing their gradual growth up to centimeter-scale sizes. During their volume increase, they push the mud out from the container resulting in horizontal and vertical propagation of the mud over cm-scales. In those experiments where the mud bulge freezes due to the evaporative cooling, the internal structure is kept in (or beneath) the icy crust. Our experimental approach reveals that mud with identical characteristics features different morphologies depending if its extrusion occurs on Earth or Mars pressure conditions.

Figure 4. Results of volumetric change measurements.

As the surface gravity on Mars is nearly three times smaller than that on Earth, we performed numerical calculations to reveal to which depth the mud undergoes boiling. We reveal that the depth of boiling would be nearly three times larger on Mars, reaching meters for sufficiently warm muds. As a consequence, the boiling and associated volumetric changes observed during our small-scale experiments may apply to meter-scales for Mars natural conditions . This suggests that the observed mounds and knobs associated with putative martian sedimentary volcanoes might indeed be related to mud volumetric changes in response to surface exposure. We also show that mud flows on Mars, and elsewhere in the Solar System, could behave differently to those found on Earth since mud dynamics are affected by the formation of bubbles in response to the different atmospheric pressures.

Figure 5. Simplified concept of gas migration and mud inflation due to the pressure drop.

Results of these experiments also suggest that other types of liquid, that are unstable in the low pressure environment, might behave similarly if their viscosity is high enough to prevent the bubble escape. The results presented herein also have implications for cryovolcanic phenomena on icy moons (e.g., Enceladus or Europa) or dwarf planets like Ceres.

References: [1] Komatsu et al., 2016. Icarus, 268; [2] Brož et al., 2019, JGR: Planets, 124(3); [3] Mazzini and Etiope, 2017, Earth-Science Reviews, 168. [4] Brož et al., 2020a, Nature Geoscience, 13(6); [5] Brož et al., 2020b, EPSL, 545; [6] Thielicke and Stamhuis, 2014, Journal of Open Research Software, 2(1).

How to cite: Broz, P., Krýza, O., Patočka, V., Pěnkavová, V., Conway, S., Mazzini, A., Hauber, E., Sylvest, M., and Patel, M.: Volumetric Changes of Mud on Mars: Evidence from Laboratory Simulations, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-277, https://doi.org/10.5194/epsc2024-277, 2024.

Acknowledgments. We acknowledge the support from the Agence Nationale de la Recherche (ANR, France) under the contract ANR-20-CE46-0013 entitled “PaleoSilica”, the Centre National d’Études Spatiales (CNES, France), the Centre National de la Recherche Scientifique (CNRS, France), and the French government under the France 2030 investment plan, as part of the Initiative d'Excellence d'Aix-Marseille Université - A*MIDEX AMX-21-RID-O47.

 References. ^[1]Lockwood et al. (1992) 23^rdLPSC. ^[2]Di Pietro et al. (2021) Icarus. ^[3]Frey & Jarosewich (1982) JGR Solid Earth. ^[4]Souček et al. (2015) EPSL. ^[5]Salvatore & Christensen (2014) JGR Planets. ^[6]Costard et al. (2017) JGR Planets. ^[7]Costard et al. (2019) JGR Planets. ^[8]Oelher & Allen (2010) Icarus. ^[9]Carter et al. (2023) 54^thLPSC. ^[10]Pineau et al. (2020) Icarus. ^[11]Hemmi & Miyamoto (2018) Geosciences. ^[12]Broz et al. (2023) Earth Surf. Dyn. ^[13]Kreslavsky & Head (2002) JGR Planets.

 Figure 1. A. to D. PHS (green) and HySi (cyan) CRISM detections in four sites within the Acidalia’s TT over CTX background, North is up. For each, coordinates of CRISM cubes are provided. Scale bars indicate 2000m.

How to cite: Pineau, M., Carter, J., Lagain, A., Mangold, N., Le Deit, L., and Quantin-Nataf, C.: Mineralogical Evidence of Sedimentary Volcanism for the Formation of the Thumbprint Terrains in Acidalia Planitia, Mars., Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-906, https://doi.org/10.5194/epsc2024-906, 2024.

On Earth, the characteristics of fluvial erosion (drainage direction, discharge, drainage area, and channel complexity) depends on climate (rain fall) and tectonic development (uplift rate, slope, and relief). Mars is a planet that experienced erosion driven by liquid water in its history, but its tectonics are vastly different from Earth’s. The formation of the Tharsis dome, a vast volcanic province represents a major magmato-tectonic upheaval for the planet. This event occurred during the early history of Mars, when it had an active water cycle, evidenced by the formation of fluvial valleys that have been preserved until present. The timing of the rise of the Tharsis province and the carving of valley network are debated (Phillips et al., 2001; Bouley et al., 2016). If they simultaneously occurred, it is likely we would see the tectonic change in Mars’ surface due to Tharsis affect the characteristics of fluvial erosion as listed above. Conversely, if both events are separate in time then the valley networks will have formed on a strictly Pre-Tharsis Mars, or a Post-Tharsis Mars. Nonetheless it is important to settle this debate in order to determine precisely the surface conditions prevailing during a period of time when the red planet could have seen the development of life.

This study is the first to model the influence of the growth of Tharsis on the formation of the Valley networks. We present results from numerical simulations performed using a Stream Power Law (SPL) algorithm on Mars during the emplacement of Tharsis to assess how the patterns of drainage are affected by the tectonic development of the planet, in an attempt to reproduce the observed distribution of valley networks and their geometry. We find that the tectonic change due to Tharsis measurably affects drainage systems across the planet and adjust the timing until we reach a close match between our modelled network and the observed one. As Tharsis emplaces and regions uplift and subside in response, their local slope and relief is altered, having the effect of dispersing networks, lowering discharge, drainage area and erosion rate in uplifting regions and converging networks, increasing discharge, coalescing drainage areas, and increasing erosion rate in subsiding regions. This causes a preservation of older networks in the uplifting regions and a replacement of networks in subsiding regions, making it appear that certain valley networks are older and others are younger (Fig. 1). This suggests that not only is Tharsis emplacement necessary to produce the relief needed to form the observed networks on Mars today, but that a prolonged (and likely intermittent) period of fluvial activity on Mars must have occurred concurrently with the emplacement of Tharsis. In conclusion, Mars offers the opportunity to understand how landscape may evolve under a wider range of tectonic regimes, compared to Earth.

Fig 1. A map of Mars with total topographic change hypothesised to have been caused by Tharsis (Matsumaya and Manga 2010), overlain with SPL model results divided into 15°x15° bins and coloured based on when in the model time the valley networks best match with observations when the two are compared using a Discrete Fourier Transform (DFT) signal comparison. Where the observational data best matches model data Pre-Tharsis, at 3.8 Ga model time is marked in black, Mid-Tharsis at 3.65 Ga in dark gray, and Post-Tharsis at 3.5 Ga in light gray. Crosses on the map denote valley networks dated by Hoke & Hynek (2009) with their respective ages.

References

Bouley, S., Baratoux, D., Matsuyama, I., Forget, F., Séjourné, A., Turbet, M., and Costard, F.: Late Tharsis formation and implications for early Mars, Nature, 531, 344–347, https://doi.org/10.1038/nature17171, 2016.

Hoke, M. R. T. and Hynek, B. M.: Roaming zones of precipitation on ancient Mars as recorded in valley networks, Journal of Geophysical Research: Planets, 114, https://doi.org/10.1029/2008JE003247, 2009.

Matsuyama, I. and Manga, M.: Mars without the equilibrium rotational figure, Tharsis, and the remnant rotational figure, Journal of Geophysical Research: Planets, 115, https://doi.org/10.1029/2010JE003686, 2010.

Phillips, R. J., Zuber, M. T., Solomon, S. C., Golombek, M. P., Jakosky, B. M., Banerdt, W. B., Smith, D. E., Williams, R. M. E., Hynek, B. M., Aharonson, O., and Hauck II, S. A.: Ancient Geodynamics and Global-Scale Hydrology on Mars, Science, 291, 2587–2591, https://doi.org/10.1126/science.1058701, 2001.

How to cite: Davies, H. S., Bouley, S., Baratoux, D., and Braun, J.: What role did Tharsis formation have on the geomorphological history of Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-503, https://doi.org/10.5194/epsc2024-503, 2024.

Tyrrhena Terra, a region located in the cratered Noachian highlands between Hellas and Isidis Planitia on Mars, is distinguished by its extensive presence of hydrated minerals. As a key region for understanding the early Martian aqueous environment, various evidence for hydrated alteration has been reported by previous studies (Loizeau et al., 2012; Rogers, 2011; Carter et al., 2013; Bultel et al, 2015; Viviano et al., 2023). However, the effects of varying crater scales on mineral formation and distribution have not yet been well understood.

In this comprehensive investigation, we identified hydrated minerals across Tyrrhena Terra (2°N-22°S, 64°E-104°E), using near-infrared hyperspectral data from the Compact Reconnaissance Imaging Spectrometer for Mars (CRISM). Along with high-resolution imaging from the Context Camera (CTX) and the High-Resolution Imaging Science Experiment (HiRISE), we explored the relationship between impact crater size and degradation rate and the diversity of hydrated minerals. Based on crater diameter, the observations were systematically categorized into five groups: away from craters, <5 km, 5-10 km, 10-20 km, and >20 km. Additionally, craters larger than 20 km were subdivided based on degradation morphology into three types (Mangold et al., 2012): type I (highly degraded), type II, and type III (least degraded).

Our findings showed a widespread distribution of phyllosilicates, including Fe/Mg-smectites, chlorite, kaolinite, mica, and their mixtures, as well as other hydrated minerals such as prehnite, zeolite, carbonate, and hydrated silica (Figure 1). We observed a significant increase in mineral diversity with larger crater diameters, even when excluding minerals located on central peaks (Figure 2). Such trend indicates that the subsurface compositions vary significantly with depth, since the post-date impact generated hydrothermal deposit mostly found in the central peak of craters. By excluding minerals on central peaks, we reduced the influence from impact-driven hydrothermalism for hydrated minerals formation, thus maximized the role of impact events to be the excavation of subsurface minerals.

Figure 1. The distribution of hydrated minerals in Tyrrhena Terra.

Figure 2. The variation of the hydrated minerals with the diameter of impact crater.

Notably, for craters larger than 20 km, the mineral diversity also shows a strong correlation with the degradation rate of craters (Figure 3). Only Fe/Mg-smectites and chlorite were observed in highly degraded type I craters, likely dating back to the Noachian period, while prehnite, zeolite, carbonate, and hydrated silica were more common in less degraded type II and type III craters. Moreover, the context of Fe/Mg-smectites exposures in Tyrrhena Terra is particularly intriguing. Fe/Mg-smectites were prevalent not only in type I craters, but also in ancient sediments that are incised by valley networks, isolated from impacts, and with small amount of Al-smectite mixtures observed (Figure 4c); as well as in exposures of small impacts (less than 5 km in diameter) on ancient sediments (Figure 4b, CRISM image FRT00014005 and FRT00024086). In contrast, the interiors of large, well-preserved craters exhibit a complex diversity of minerals, including those formed under relatively high temperatures, such as prehnite and zeolite. It appears that minerals formation involves processes beyond only impact events. Such difference revealed clear evidence for active surface alteration in Noachian period, resulting in the transformation of hydrated minerals.

Figure 3. The variation of the hydrated minerals with the degradation rate of impact crater larger than 20 km.

Figure 4. Key examples for the context of Fe/Mg-smectites and Al-smectites (mixed with Fe/Mg-smectite) in Tyrrhena Terra. (a) CRISM spectral ratios (top and middle frame) compared to Fe/Mg-smectites and Al-smectites laboratory spectra from USGS and RELAB (bottom frame). (b) Fe/Mg-smectites exposed by small craters on ancient sediments. (c) Fe/Mg-smectites on ancient sediments and isolated from impacts (in magenta), as well as small exposures of Al-smectites mixed with Fe/Mg-smectites (in cyan, extracted from Carter et al., 2023). The RGB channels are Band 233, Band 78, and Band 13 respectively for all CRISM images.

Our findings highlight a regional source-to-sink geological process where minerals were transformed and redistributed. These minerals are excavated by local impacts, subsequently altered by surface weathering processes, and transported through fluvial activity to their present locations. This suggests that more than one period of hydrated alteration was active in early Martian history, likely predating the formation of widespread valley networks, with each period lasting long enough for hydrated minerals to be formed or altered. This research offers new insights into the early Martian aqueous history and its past environmental conditions.

References:

Bultel, B., Quantin-Nataf, C., Andréani, M., Clénet, H., & Lozac’h, L. (2015). Deep alteration between Hellas and Isidis basins. Icarus, 260, 141–160.
Carter, J., Poulet, F., Bibring, J.-P., Mangold, N., & Murchie, S. (2013). Hydrous minerals on Mars as seen by the CRISM and OMEGA imaging spectrometers: Updated global view. Journal of Geophysical Research: Planets, 118(4), 831–858.
Carter, J., Riu, L., Poulet, F., Bibring, J.-P., Langevin, Y., & Gondet, B. (2023). A Mars orbital catalog of aqueous alteration signatures (MOCAAS). Icarus, 389, 115164.
Loizeau, D., Carter, J., Bouley, S., Mangold, N., Poulet, F., Bibring, J.-P., Costard, F., Langevin, Y., Gondet, B., & Murchie, S. (2012). Characterization of hydrated silicate-bearing outcrops in Tyrrhena Terra, Mars: Implications to the alteration history of Mars. Icarus, 219(1), 476–497.
Mangold, N., Adeli, S., Conway, S., Ansan, V., & Langlais, B. (2012). A chronology of early Mars climatic evolution from impact crater degradation. Journal of Geophysical Research: Planets, 117(E4).
Rogers, A. D. (2011). Crustal compositions exposed by impact craters in the Tyrrhena Terra region of Mars: Considerations for Noachian environments. Earth and Planetary Science Letters, 301(1-2), 353–364
Viviano, C. E., Beck, A. W., Murchie, S. L., Dapremont, A. M., & Seelos, F. P. (2023). Heterogeneity of the Noachian crust of Mars using CRISM multispectral mapping data. Geophysical Research Letters, 50(5), e2022GL102711.

How to cite: Wu, Y., Mangold, N., Liu, Y., Carter, J., Wu, X., Pan, L., Huang, Q., and Zou, Y.: Comprehensive analysis of the alteration of Tyrrhena Terra: Implications for source-to-sink processes on Mars, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-364, https://doi.org/10.5194/epsc2024-364, 2024.

Sulfate minerals are an integral component of the martian surface and understanding the formation and alteration of these minerals provides clues about their geochemical environment. One sulfate phase in particular has been intriguing Mars scientists for over 15 years. An unusual spectral band at 2.236 µm was discovered in CRISM spectra of Mars at the plateau bordering Juventae Chasma [1] and in Aram Chaos [2]. This spectral band does not line up with any known minerals, but is observed for FeSO₄OH, a new mineral formed by heating hydrated iron sulfates [3]. Crystal structure diagrams indicate that FeSO₄OH has a structure similar to that of szomolnokite (FeSO₄•H₂O), but with OH replacing H₂O [4]. Experiments heating szomolnokite, rozenite, and melanterite in the lab resulted in production of some FeSO₄OH at 150 °C from rozenite and melanterite after 30 minutes and complete transformation at 200 °C after 30 minutes, while reaction of szomolnokite required longer heating at 200 °C and/or elevated temperatures for transformation to FeSO₄OH. Additional lab experiments also demonstrate that oxygen is required for reaction of these hydrated ferrous sulfates to form FeSO₄OH [4].

Application of FeSO₄OH spectra to Mars has benefitted from improved processing techniques [5] and mapping algorithms [6] for CRISM images that has enabled characterization of smaller spot sizes with cleaner spectra. Exposures of this unusual ferric sulfate phase with spectral features near 2.23 µm at Aram Chaos closely resemble pure FeSO₄OH formed in the lab, while the thinner units on the Juventae plateau are either mixed with other components or represent incompletely formed FeSO₄OH phases. This Fe hydroxysulfate is currently associated with monohydrated sulfate (MHS) outcrops at Aram Chaos, although polyhydrated sulfate (PHS) outcrops are also present nearby. In contrast, only PHS outcrops are currently observed on the Juventae plateau.

CRISM spectra with a spectral band at 2.225-2.238 µm were observed at several locations at Aram Chaos. The spectra of these units also contain accompanying features at 1.48, 1.82, 2.19, and 2.37 µm (Fig. 1). Small variations in the ~2.23 µm band are attributed to changes in the Fe-Mg chemistry. Pure FeSO₄OH has a band at 2.236 µm, while heated FeMg-MHS has a band at 2.226 µm. The FeSO₄OH units at Aram Chaos are typically found adjacent to MHS outcrops (Fig. 2), including both Fe-rich MHS (similar to szomolnokite) with bands near 2.11-2.12 and 2.40 µm and kieserite (MgSO₄•H₂0) with spectral bands near 2.14 and 2.41 µm. Spectra of szomolnokite and kieserite measured at colder, Mars-like temperatures have bands at ~2.11 and 2.14 µm [7], similar to these observations. The MHS spectral units more consistent with kieserite are darker than the szomolnokite-like MHS units and are covered by debris and ripples (Fig. 2C). Some of the MHS units at Aram Chaos may have formerly been mixtures of szomolnokite and kieserite, where the szomolnokite transformed to FeSO₄OH and the kieserite remained. Alternatively, polyhydrated Fe and Mg sulfates may have been present that altered to form FeSO₄OH, szomolnokite, and kieserite, depending on variations in the geochemistry. The reaction to form FeSO₄OH proceeds as Fe²⁺ in szomolnokite is oxidized to Fe³⁺ while the H₂O loses a proton to form OH. Some of the MHS outcrops include weak bands near 2.23 µm, indicating partial alteration to form a mixed phase containing some MHS and some FeSO4OH (Fig. 3). The PHS spectra contain bands near 1.44 and 1.93-1.95 µm and a drop in reflectance near 2.42 µm (Fig. 3), similar to spectra of rozenite (FeSO₄•4H₂0) and starkeyite (MgSO₄•4H₂0).

Spectra at the Juventae plateau were collected from thin light-toned layered deposits, including spectral signatures due to PHS and FeSO₄OH, and pyroxene-bearing units (Figs. 3-4). The stratigraphy of the outcrops shows a pyroxene-bearing substrate below the light-toned layered materials and a different pyroxene-bearing caprock unit covering the Fe sulfates (Fig. 4). Thin units containing spectral features consistent with PHS (blue) and FeSO₄OH (red) are observed, and some of these thin units include spectral features due to both materials. Morphologies of these primary four units are displayed in Fig. 5. The pyroxene bearing substrate (dark cyan) is flatter with extensive polygonal fracturing, whereas the pyroxene-bearing caprock (green) is partially covered by ripples and appears hilly and uneven in topography due to differential erosion. The textures of the PHS- and FeSO₄OH-bearing units are distinct from those of the pyroxene-bearing units, but appear related to each other with fine-scale layering that varies in brightness, color, and fracturing.

The presence of highly pure FeSO₄OH outcrops neighboring MHS at Aram Chaos and less pure outcrops of FeSO₄OH neighboring PHS on the plateau NW of Juventae Chasma indicate an active geochemical history in Mars’ past. The hydrated sulfates likely formed in evaporative environments, while the FeSO₄OH likely formed through heating. The FeSO₄OH-bearing units at the Juventaue plateau could be mixed with spectrally neutral components that dilute the FeSO₄OH spectral features. Coordinated characterization of the near-infrared (NIR) and mid-IR spectral features of Fe sulfates (Fig. 6) is enabling a better understanding of the spectral features due to FeSO₄OH in these intriguing outcrops on Mars.

Acknowledgements: The authors are grateful for support from NASA MDAP #80NSSC21K1103, NASA SSW #80NSSC23K0032, Austrian Science Fund FWF #P34227-N, and Europlanet Transnational Access funds.

References: [1] Bishop J.L. et al. (2009) Mineralogy of Juventae Chasma…, JGR, 114, doi:10.1029/2009JE003352. [2] Lichtenberg K. A. et al. (2010) Stratigraphy of hydrated sulfates in the sedimentary deposits of Aram Chaos, Mars, JGR, 115, doi:10.1029/2009JE003353. [3] Bishop J. L. et al. (2024) Characterizing the spectral properties of a new FeSO4OH phase observed on Mars, LPSC, #1880. [4] Meusburger J. M. et al. (2024) Ferric hydroxysulfate on Mars and its formation from ferrous sulfate hydrates, 10^th Mars Conference. [5] Itoh Y. & M. Parente (2021) A new method for atmospheric correction and de-noising of CRISM data, Icarus, 354, 114024. [6] Saranathan A.M. & M. Parente (2021) Adversarial feature learning for improved mineral mapping of CRISM data, Icarus, 355, 114107. [7] Yeşilbaş M. et al. (2024) Low-temperature reflectance spectra of szomolnokite and applications for their detection on Mars, LPSC, #2035.

How to cite: Bishop, J. L., Meusburger, J. M., Talla, D., Weitz, C. M., Parente, M., Gross, C., Saranathan, A. M., Itoh, Y., Gruendler, M. R. D., Howells, A. E. G., Yeşilbaş, M., Hiroi, T., Schmitt, B., Maturilli, A., Al-Samir, M., Bristow, T. F., Lafuente, B., and Wildner, M.: Characterizing the New Mineral FeSO4OH on Mars and Describing its Geochemical History and Association with Other Sulfates , Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-289, https://doi.org/10.5194/epsc2024-289, 2024.

Introduction

The Emirates Mars InfraRed Spectrometer (EMIRS) instrument onboard the Emirates Mars Mission (EMM) “Hope” probe is a Fourier Transform Infrared spectrometer that has been observing the Martian surface and atmosphere between 6 and 100 μm since February 2021 [1, 2]. The unique orbit of EMM allows EMIRS to observe the entire Martian disk at each observation, covering all the surface of the planet in ~4 orbits, which corresponds to ~5° of L_s, or 10 Earth days.

[3] used the surface temperature retrievals from EMIRS to detect and monitor the presence of H₂O and CO₂ ice on the surface of the planet, including the temporal and local time evolution of the CO₂ frost that can appear at Martian equatorial latitudes. This frost has been observed in the second half of the night around the equinoxes. Small crystals and optically thin layers are expected from the work of [4] using THEMIS data. However, as EMIRS provides fully resolved spectra whereas THEMIS is a multi-band instrument, we will be able to strengthen the constraints on the physical properties of these deposits. And monitor for the first time the evolution of the ice properties during its condensation phase from 12 a.m. to 6 a.m. thanks to the unique temporal coverage of EMM instruments.

Data & Methods

Based on the first identification of the EMIRS pixels mostly covered by CO₂ ice presented in [3], we use here the full spectroscopic power of the instrument to characterize and constrain the physical properties of these icy deposits (crystal size, thickness) and their temporal evolution.

First, we identify areas where CO₂ surface frost has been detected at different local times over a few degrees of L_s(to have a sufficient spatial and temporal coverage), then we bin the data over temporal bins of 1-hr and compare the averaged spectra for each bin. The spatial extent of the EMIRS pixel footprints is computed using the SPiP Python module [5].

Results

First, we observe that CO₂ ice spectral signatures are observed in the EMIRS spectra when predicted by the temperature criterion, between midnight and 6 a.m., but mostly from 3 a.m. Comparison of spectra presented in Figure 1a with models from [4] may suggest optically thin layers of CO₂ ice, with a thickness of ~ 100 μm, when models can predict condensation of up to a few tens of microns of ice just before sunrise.

As the surface temperature remains below the freezing point of CO₂ during the second half of the night, and CO₂ is always available at the surface from the atmosphere, condensation is expected to occur over the night hours. However, no significant temporal variations of the spectra have been observed over the night. The differences between spectra that can be seen in Figure 1 are mostly associated with spatial variability of the surface emissivity, as can be seen on panels c & d. This is not in favor of the scenario of an accumulation of ice over the night which would increase progressively the frost thickness and/or the size of the crystals. Thus, the lack of variability over the night may suggest that the condensation may occur within the surface regolith and/or the subsurface, as suggested by [6].

Figure 1: EMIRS spectra from EMM orbits 171 to 180 where surface CO₂ frost has been detected between 150°W and 90°W and between 0°N and 60°N (North of Tharsis), averaged over temporal bins of 1-hr (a) and footprints of the considered EMIRS pixels (b). The colors of the spectra and pixels indicate the local time range for the observations. (c & d) EMIRS average spectrum and pixel footprints over the same region for observations between 11 a.m. and 1 p.m. to provide a daytime emissivity reference to be compared with the nighttime spectra.

Conclusion & Perspectives

In this work, we present the first spectroscopic monitoring of surface nighttime CO₂ frost under equatorial latitudes as a function of the local time. We do not observe significant variation of the spectra over the night but a spatial variability is present. This suggests that the frost condensation likely occurs within the regolith and/or close subsurface rather than a simple accumulation of ice at the surface overnight. Further spectral modeling of the emissivity will be needed to assess the physical properties of these deposits considering different scenarios including notably: “dirty” CO₂ ice, subsurface CO₂ ice, or slab of CO₂ ice.

Acknowledgments

This work was funded by the Emirates Mars Mission project under the Emirates Mars Infrared Spectrometer instrument via The United Arab Emirates Space Agency (UAESA) and the Mohammed Bin Rashid Space Centre (MBRSC).

A. S. also acknowledges funding by CNES.

References

[1] Amiri, H. E. S. et al. (2022), SSR, 218, 4. [2] Edwards, C. S. et al. (2021), SSR, 217, 77. [3] Stcherbinine, A. et al. (2023), GRL, 50, e2023GL103629. [4] Piqueux, S. et al. (2016), JGR: Planets, 121, 1174-1189. [5] Stcherbinine, A. (2023), Zenodo, doi:10.5281/zenodo.7714204. [6] Lange, L. et al. (2022), JGR: Planets, 127, e2021JE006988.

How to cite: Stcherbinine, A., Edwards, C., Haberle, C., Smith, M., Lange, L., and Pilorget, C.: Spectroscopic Characterization and Evolution of Martian Nighttime CO2 Frost at Equatorial Latitudes with EMM/EMIRS, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-50, https://doi.org/10.5194/epsc2024-50, 2024.

INTRODUCTION

Dome-shaped uplifted and fractured terrain has been observed at the surface of volcanic regions on Mars, amongst other terrestrial planets and moons in our Solar System. Such dome-shaped ground deformation features include, for example, domes in Elysium Planitia and floor-fractured craters (e.g., Farrand et al., 2011; Jozwiak et al., 2012). These features are inferred to have formed by the emplacement and inflation of sill- and laccolith-shaped magma bodies in the shallowest 1-2 km of the crust (e.g., Michaut et al. 2013). This interpretation is informed by monitored magma intrusion events and geological observations at eroded outcrops of volcanic plumbing systems on Earth. On Mars and other planetary bodies, however, only the final surface deformation features can be observed from orbit. Besides studying analogue outcrops on Earth, analytical and numerical models are used instead to understand the emplacement dynamics and the deformation of the planetary crust´s rocks. A mismatch exists, however, between the oversimplified intrusion geometry and the linearly elastic response to magma intrusion assumed by most numerical models, and the complex intrusion geometries and mechanical response of host rocks to magma-induced stresses observed on Earth in exposed volcanic plumbing systems. Strain can accumulate along large-scale discontinuities in the overburden rocks, making the investigation of the emplacement mechanisms by traditional continuum models difficult. To investigate the ill-understood effect on magma-induced deformation of the lower gravity on Mars, and the Moon, due to their smaller mass compared to Earth, we compare simulations of laccolith inflation in the Discrete Element Method (DEM) under gravity of the Moon, Mars and Earth.

METHOD

We used the two-dimensional (2D) DEM particle flow code PFC2D (Itasca Consulting Group, Inc.) to simulate the inflation of a half-ellipsoid laccolith in a rock medium represented as a particle-based assemblage. Previously, we have shown how this model can indicate fracturing and highly discontinuous deformation, as well as visualise the localization of subsurface strain and corresponding deformation (for details see Morand et al., 2024). To simulate laccolith inflation, the laccolith-shaped pressure source is inflated by increasing the area of the 2D particles (Figure 1). We ran inflation simulations in rocks of different toughness and stiffness representing a range of realistic crustal strengths, under the gravity of the Moon (g = 1.62 m.s^-2), Mars (g = 3.71 m.s^-2), and Earth (g = 9.81 m.s^-2).

Figure 1. General view of the DEM model in PFC2D with rock particles (grey) and the initial particle-based, laccolith-shaped magmatic body (red) at 1 km depth. The subsets display the magma reservoir area in its initial state before inflation and in its final state after 25% of inflation.

RESULTS AND DISCUSSION

For equal rock stiffness and amounts of intruded magma, our model results show that we can expect the same volume of magma injected at the same depth to induce more vertical surface displacement on Mars, due to the lower gravity there compared to Earth. Nevertheless, rock toughness and rock stiffness control the amount of fracturing more than gravity does. These findings are even better expressed under Lunar gravity. Our finding implies that both gravity and realistic crustal rock strength are essential parameters to account for in modelling efforts of magma-induced deformation on moons and planets such as Mars with a size and mass that is significantly different from that of Earth.

Our model results will induce a better understanding of the emplacement and architecture of shallow magmatic intrusions below magma-induced uplifted terrain and floor-fractured craters on Mars. Verification of this new modelling application with detailed structural information from analogue magma intrusions on Earth, such as Permian trachy-andesite intrusions in the Intra-Sudetic Synclinorium in SW Poland, allow us to investigate the characteristics of the cryptic magma bodies underlying surface doming on Mars and the Moon, while simultaneously improving the interpretation of volcanic unrest signals on Earth.

REFERENCES

Farrand, W.H. et al. (2011) “Spectral evidence of volcanic cryptodomes on the northern plains of Mars.” Icarus, 211(1), 139–156. https://doi.org/10.1016/j.icarus.2010.09.006

Jozwiak, L.M. et al. (2012) “Lunar floor-fractured craters: Classification, distribution, origin and implications for magmatism and shallow crustal structure.” Journal of Geophysical Research E: Planets, 117(11), 1–23. https://doi.org/10.1029/2012JE004134

Michaut, C., et al. (2013) “Magmatic intrusions and deglaciation at mid-latitude in the northern plains of Mars.” Icarus, 225(1), 602-613. https://doi.org/10.1016/j.icarus.2013.04.015

Morand, A. et al. (2024) “Fracturing and dome-shaped surface displacements above laccolith intrusions: Insights from Discrete Element Method modeling.” Journal of Geophysical Research: Solid Earth, 129, e2023JB027423. https://doi.org/10.1029/2023JB027423

How to cite: Poppe, S., Cornillon, A., Morand, A., Harnett, C. E., and Mège, D.: Displacement and fracturing around laccolith intrusions affected by gravity on Mars versus the Moon, and Earth, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-761, https://doi.org/10.5194/epsc2024-761, 2024.

Introduction

Sublimation experiments performed in vacuum chambers have proven useful for understanding the processes that occur on the surfaces of volatile-rich celestial bodies without significant atmospheres [1, 2, 3]. To date, these experiments have focused on the sublimation of volatile ices without considering the possibility of a liquid phase being present. The most prominent example is Mars, the surface of which appears as an ultra-arid cold desert nowadays, but the presence of liquid water in its past is essentially assured [4, 5].

On Mars, impact events or seasonally increased insolation might locally thaw near-surface permafrost and release liquid water, even long after transitioning from a more humid to an ultra-arid and low-pressure environment. The liquid water could solute and transport salts away from one location to another, where they could form deposits after the eventual evaporation of the water.

Suppose salt crusts or related morphological features of the Martian surface can be reproduced in sublimation chambers using saline solutions and regolith analogues. This would shed light on how salt deposits on Mars formed after the planet lost most of its atmosphere.

Scientific objectives are:

- What physical environmental parameters and chemical composition affect the sustainability of surface-near aqueous solutions in a cold, low-pressure environment?

- Does liquid water sublimate from the sample surface and form salty crusts or efflorescences?

- Do saline ices and brines form near-surface features such as ice lenses?

Methods

In a vacuum sublimation chamber, environmental conditions on the Martian surface were reproduced, and the conditions under which saline solutions can migrate through a regolith analogue were investigated. The focus was on the parameters of salt composition and concentration, as well as the grain size of the regolith analogue. In particular, Mg²⁺, Fe³⁺, Cl^-, and SO₄^2- based salts were used since they are observed in specific areas of the Martian surface [6, 7].

Mars regolith analogues were saturated with brines and placed in a transparent sample container in a sublimation chamber. The environmental conditions in the chamber were adapted to the Martian summer and account for 243 K temperature, 5 mbar atmospheric pressure, and 600 Wm^-2 irradiance.

Each experiment lasted several days, during which the lateral and horizontal sample surfaces were permanently monitored with cameras. After completing the sublimation experiments, the sample's internal layer structure was analysed. The reflectance of each sample surface was measured in the spectral range of 400-2500 µm. Eventually, the cone tip penetration method was used to test the resistance of the surface to mechanical stress.

We carried out experiments on two scenarios with different salt concentrations, in which the grain size of the regolith analogue was varied at low and high salinity. The used grain sizes were 212-500 µm, 1000-2000 µm, and 212-2000 µm.

Results

After the sublimation experiments, the samples generally presented four different layers (Fig. 1).

- A solidified crust formed on the surface. The thickness of the crust was determined by the grain size, whereby the thicknesses increased with increasing grain sizes. With increasing grain size, the crust's cone tip resistance (CTR) was reduced by an order of magnitude, respectively. However, the resistance to mechanical stress also depended strongly on the salinity. For the same grain size, the CTR for the high salinity scenario was about one magnitude higher than in the low salinity scenario. Analyses of the reflectance spectra showed a clear blue slope of salt-rich surfaces, which became more distinct with decreasing regolith grain size (Fig. 2). However, no pure salt deposits could be produced on the surface, even in the high salinity scenario.

- Below the crust, a layer of dry, unconsolidated material formed. Notably, in the samples of the low salinity scenario, salt crystallised and gave the layer a bright appearance.

- Then, a slightly moist layer followed, and dark-yellowish crystals formed. The thickness of the layer was comparable for both salinities and only increased slightly with increasing grain size.

- The bottom layer was moist with brines and contained considerable quantities of water ice crystals. The depth at which this layer formed was determined by grain size, not salinity. Small grain sizes favoured the formation of larger ice crystals, which grew to 5 mm in size. The ice crystals were clear and had a lower salt content compared to the initial solution.

Fig 1: Schematic representation of the resulting layers after completion of the sublimation experiments. The camera images show examples of the layers with various grain sizes and salinities.

Fig 2: Reflectance spectra of sample surfaces with a regolith grain size of 212-500 µm. With increasing salinity of the sample, the spectra show a distinct blue slope.

References

[1] Poch, O., et al. 2016, Icarus, 267, 154-173. doi:10.1016/j.icarus.2015.12.017

[2] Mc Keown, L.E., et al. 2017, Sci Rep 7, 14181. doi:10.1038/s41598-017-14132-2

[3] Haack, D., et al. 2021, Astronomy & Astrophysics, 649, A35. doi:10.1051/0004-6361/202140435

[4] Irwin III, R.P., et al. 2004, Journal of Geophysical Research: Planets, 109(E12). doi:10.1029/2004JE002287

[5] Wray, J., 2021, Annual Review, 49, 141-171. doi:10.1146/annurev-earth-072420-071823

[6] Mangold, N., et al. 2008, Icarus, 194, 519-543. doi:10.1016/j.icarus.2007.10.021

[7] Osterloo, M.M., et al. 2008, Science, 319(5870), 1651–1654. doi:10.1126/science.1150690

How to cite: Haack, D., Kaufmann, E., and Hagermann, A.: Evolution of brines in regolith analogue under simulated Martian surface conditions, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-85, https://doi.org/10.5194/epsc2024-85, 2024.