Post contributed by Alistair Blance, The Open University, UK
During an impact on Mercury’s surface, material is ejected from the forming impact crater. As Mercury has only a tenuous atmosphere, ejected material travels predominantly ballistically, creating an ejecta deposit around the crater that thins gradually with increasing distance. However, large deposits emplaced by ground-hugging flows can be found around some impact craters on Mercury (Image 1). Evidence for flow includes material being diverted around obstacles, a steep edge or distal ridge at deposit margins, and a lobate shape to several examples. Some flow deposits extend outwards around a whole crater, but often they are confined within topographic lows adjacent to the crater. To help assess the origin of these features, it is useful to compare them to similar features across the Solar System. This comparison may also indicate how differences between the planets can influence the development of flows around craters.
Image 1: Flow deposits around craters on Mercury. Deposit boundaries indicated with red triangles. (A) Flow deposit extending from the central crater into an underlying crater in the top right of the image. Steep margins with a lobate shape suggest emplacement by flow. Image taken from MESSENGER MDIS BDR Global Basemap. (B) A crater with two sections of flow deposit extending into the underlying crater in the bottom right of the image. Image taken from MESSENGER MDIS frame EW0260906588G. (C) Sketch map of the image in B. Shows the two sections of flow deposit in red, with hypothesised direction of emplacement shown with red arrows. The deposit appears to have been diverted around a central peak within the underlying crater, providing evidence for emplacement via ground-hugging flow.
Post contributed by Dr Anna Grau Galofre, Laboratoire de Planetologie et Geosciences, CNRS UMR 6112, Nantes Universite, Universite d’Angers, Le Mans Universite, France.
Glaciers, ice sheets, and ice masses in general, are both agents of landscape evolution and a part of the landscape itself, when considered at human time scales, with surfaces that evolve and record changes in environmental conditions. The ages of glaciers and ice sheets on Earth, which range from thousands to millions of years, dwarf in comparison with the ages of debris-covered ice masses on Mars. Debris-covered martian ice deposits range up to hundreds of millions of years old, perhaps approaching 1 billion years old in some locations. Image 1 shows complex patterns on the surface of a debris-covered glacier on Mars. Terrestrial glacial surface deformation patterns can record changes in environmental conditions, which are also captured in the variability in chemistry and crystal structure seen in their internal stratigraphy. The morphology of a glacier surface, for example, may reflect episodes of warmer climate, and therefore enhanced ice flow; banding patterns may reflect periodic changes in the mass balance of ice accumulation and ice loss; and crevasses may reflect episodes when ice flowed relatively fast and fractured. The surface morphology of glaciers and debris-covered ice deposits can thus help us learn about the history of climate on Earth and Mars.
Image 1. Surface banding patterns on a debris-covered, cold based ice deposit (called a ‘lobate debris apron’) flowing down a canyon wall in Reull Vallis, Mars. Extract from CaSSIS image MY34_005511_224 (footprint width is 9 km). Image credit: ESA/CaSSIS.
The polar caps of Mars are gigantic ice deposits, similar in size to Texas or France, and mostly made up of frozen H2O. Like many planets and moons, Mars is littered with impact craters – the scars of violent asteroid and comet collisions. Near to but separated from the polar caps, some craters on Mars are home to mounds of frozen H2O ice. These “ice mounds” are much smaller than the giant polar caps, but not tiny—think the size of Rhode Island or Luxembourg instead of Texas or France—and have been discovered in dozens of craters using robotic spacecraft that collect images of the surface from Martian orbit. An example of one ice mound is shown in Image 1.
Image 1. Ice mound located in Korolev Crater in the north polar region of Mars. The ice mound is about 37 miles (60 kilometers) wide. Image is a perspective view made from data returned by ESA’s Mars Express spacecraft.
Post contributed by Dr. Zach Dickeson, Department of Earth Sciences, Natural History Museum, London, UK.
At geological timescales, individual lakes are only transient features on the Earth’s surface. Conditions within lake basins change over time, eventually filling with sediment, drying up, or draining completely by steady erosion (or catastrophic breach) of the basin margin. Sites where lakes existed in the past are known as palaeolakes, and in addition to Earth, hundreds of examples have been identified on Mars (Goudge et al., 2015). Rather than a single characteristic landform, palaeolakes on Mars are typically identified by a set of topographic and morphological features such as shown in Image 1. Careful investigation of associated river valleys, deltas, layered deposits, and shoreline terraces can reveal details of a lake’s origin, development, and eventual demise (Dickeson et al., 2022).
Image 1: Deposits and landforms within a palaeolake in Arabia Terra, Mars. Two past lake levels are represented by contours defined by outlet valley elevations. The inlet valley is visible at the bottom of the image where it enters the palaeolake basin. A broad, flat-surfaced sediment fan with a steep front sits at a higher elevation, and is incised through by the inlet valley, which continues into the basin to terminate at the apex of a lower and more heavily weathered sediment fan. The two fans are interpreted as deltas, and match the two different fill levels within the palaeolake. Context Camera image D12_031993_2114_XN_31N014W; http://viewer.mars.asu.edu/viewer/ctx/D12_031993_2114_XN_31N014W; NASA/JPL/MSSS/University of Arizona.
The palaeolake shown in Image 1 is one of six recently discovered, which form a hundreds of kilometers long chain (Dickeson et al., 2022). As with the majority of palaeolakes on Mars, this site was identified where the top end of a valley intersects the edge of a topographic basin, indicating that the basin was once full of water to the point of overflowing through the valley. By tracing contours around the basin at the elevation of the valley head, two past lake levels can be discerned; a higher level existing when the lake was initially filled, and a lower level after erosion of the outlet valley. These relationships can be seen in Image 2, which shows an oblique 3D view across the basin towards the fans and inlet valley. In this example, the two fill levels defined by the outlet valley are further supported by the presence of two possible deltas (labelled as the upper and lower fans in Image 2) at an inlet valley on the opposite edge of the basin. The higher delta matches the early lake level, formed as water rushed in to fill the basin, and the lower delta matches the later lake level, following outlet valley erosion and partial drainage of the basin. In contrast to the surrounding terrains, mesas are also observed only within the defined palaeolake extent. Displaying layered sides and surfaces dipping towards the basin centre, these formations suggest prolonged deposition within the palaeolake basin.
Image 2: A simulated view from near the palaeolake outlet, looking back towards the inlet valley and deltas in the distance, and a mesa with layered deposits in the foreground. Image constructed by overlaying a 0.5 m/pixel HiRISE image (ESP_031993_2115; https://www.uahirise.org/ESP_031993_2115; NASA/JPL/University of Arizona) over a 2 m/pixel digital elevation model (https://doi.org/10.6084/m9.figshare.19071560.v1), with 2x vertical exaggeration applied.
Palaeolakes constitute some of the strongest evidence for the widespread occurrence of water on ancient Mars, often exhibiting corroborative morphological, mineralogical, and sedimentary evidence in a single locality (Alemanno et al., 2021). In the search for past water and life on Mars, palaeolakes are therefore attractive targets for exploration, and both the Curiosity and Perseverance rovers are exploring palaeolake basins at Gale (Palucis et al., 2016) and Jezero (Farley et al., 2020) craters respectively. However, these two sites only represent a portion of palaeolake diversity on Mars, and the study described here highlights the detailed findings which are possible with existing orbital datasets.
Post contributed by Dr Valentina Galluzzi – National Institute of Astrophysics, Institute for Space Astrophysics and Planetology (INAF-IAPS), Italy.
The ESA/JAXA BepiColombo spacecraft is on a long journey into orbit around Mercury. This journey includes six flybys of the planet before orbital insertion in 2025. On 23 June 2022, BepiColombo accomplished its second flyby of Mercury by approaching the planet as close as 200 km from the surface. Unfortunately, this happened when night shadows were still hiding the surface from sight. It was just 5 minutes after closest approach that the Monitoring Cameras (M-CAMs, three in total) mounted on the BepiColombo’s Mercury Transfer Module (MTM) could start taking the first snapshots of the planet with enough light. This caused the pictures to show dramatic sunrise shadows along the terminator. One of the first regions to be imaged was the Eminescu area of Mercury, as seen in Image 1.
Image 1: Composite image of the Eminescu region of Mercury made with BepiColombo/MTM M-CAM#2 “Image 02” taken during BepiColombo’s second Mercury flyby (credits: ESA) overlain by MESSENGER/MDIS enhanced colors (Denevi et al., 2018). Some spacecraft parts are visible near the frame margins. North is to the left. For scale reference, Izquierdo crater diameter is about 150 km (see Image 2).
In this picture, the shadow-enhanced morphology comes from the original M-CAM#2 frame taken at 09:24 UTC. By georeferencing the original M-CAM photo onto the Mercury’s spatial frame (Galluzzi et al., 2022) it was possible to overlay terrain color using the MESSENGER/MDIS enhanced color mosaic (Denevi et al., 2018). The composite image helps to highlight the terrain diversity in this area, from the yellowish young smooth plains to the dark blue and rougher intercrater plains.
Post contributed by Dr. Olivier S. Barnouin, The Johns Hopkins University Applied Physics Laboratory, USA
The OSIRIS-REx (Origins, Spectral Interpretation, Resource Identification, and Security–Regolith Explorer) spacecraft successfully sampled the surface of the asteroid (101955) Bennu on October 20, 2020. The sampling event revealed that Bennu’s surface was very weak and easy to displace. But this outcome was not necessarily evident when OSIRIS-REx first arrived at Bennu, given the asteroid’s very rocky surface and apparent lack of fine sand-sized surface material, or regolith (Image 1), until we began analyzing Bennu’s landscape in detail.
Image 1. Movie of high resolution (40 cm) digital terrain model of Bennu showing its rough, rubbly surface. Bennu’s diameter is ~500m. In this movie, the sun is held fixed at at 20 sub-longitude. Movie generated with Small Body Mapping Tool (sbmt.jhuapl.edu). Credit CSA/NASA/UA/APL.
On the 29th of September 2022 the Juno spacecraft, in orbit around Jupiter, made a close pass by Jupiter’s icy moon Europa (Image 1). The spacecraft approached from Europa’s night side, passed the terminator (day-night boundary), and departed on the day side, coming within ~350 km of the surface. Juno’s visible color imager, JunoCam, snapped 4 images of Europa as the spacecraft sped by at a speed of 23.6 km/sec on its way to its closest approach to Jupiter. In an elliptical polar orbit, this was the only opportunity in the mission for Juno to get close to this moon of Jupiter and the first time since the Galileo mission ended in 2003 that any spacecraft has flown so close.
Europa is categorized as an “ocean world” with a solid ice surface over a liquid water subsurface layer. Europa is crisscrossed by numerous cracks, bands, ridges and troughs (lineaments) that record the tidal stress the moon experiences arising from the gravitational pull of Jupiter and its other moons. JunoCam’s image reveals numerous pits along the terminator. The almost complete lack of craters tells us that geologically this icy surface is very young, resurfaced by lineament formation due to tidal flexing. Callanish crater, one of the few imaged by JunoCam, is the circular feature visible in the lower right of Image 1.
Post contributed by Abhinav Jindal, Department of Astronomy, Cornell University, USA.
The European Space Agency’s Rosetta mission to 67P/Churyumov–Gerasimenko (67P) provided the most detailed ever views of a comet’s surface. By tracking the comet on its journey through the inner solar system for just over two years, Rosetta observed an increasing amount of activity and surface changes as the ice-rich surface warmed up. These observations revealed that 67P had seasons. Specifically, due to its present-day obliquity and orbit, southern summer corresponds closely with perihelion. This results in orders of magnitude greater insolation at southern latitudes than in the north [Hu et al., 2017]. Accordingly, a greater number of particles are eroded and launched from southern latitudes, with a large fraction falling back onto the surface in the cold northern hemisphere [Thomas et al., 2015; Keller et al., 2017]. This creates a hemispheric dichotomy in surface morphologies – layered bedrock units that represent the exposed nucleus are dominant in the south (“rough” terrains), while topographically smooth, sediment-covered regions dominate the north (“smooth” terrains) [Birch et al., 2017]. It was within the smooth terrains that we observed the majority of changes on 67P [Groussin et al., 2015; El-Maarry et al., 2017; Hu et al., 2017; Pajola et al., 2017a; Birch et al., 2019], making them critical to understanding the overall evolution of the comet. Our work [Jindal et al., 2022] tracked some of the most dramatic large-scale changes within these sedimentary deposits (Image 1). Specifically, in the largest and most southern smooth terrain deposit – termed Imhotep – we observed dozens of pits, bounded by multi-meter high scarps, grow and migrate across the region for many months.
Image 1: Activity in Imhotep. Animation showing georeferenced images of the Imhotep region from 05/24/2015 – 01/23/2016. Activity in the region persists from 06/05/2015 – 12/06/2015 and can be seen in the form of scarps (arcuate depressions) migrating through the region. Images were downloaded and are freely available on ESA’s Planetary Science Archive (https://archives.esac.esa.int/psa) and NASA’s Planetary Science Data System (https://pds.nasa.gov), and were georeferenced using the shapeViewer software (www.comet-toolbox.com).
Post contributed by Dr Alex Morgan, Planetary Science Institute, USA
Alluvial fans are sedimentary deposits that form at the base of mountain fronts as a channel exits a steep, confined valley onto more gently sloped and unconfined terrain (Image 1). The channel loses its ability to transport sediment due to the reduction in slope and the lateral expansion of flow, and over time sediment is deposited into a semi-conical shape. The effects of water on a landscape are the most unambiguous markers of past climatic environment, and the distribution of alluvial fans provides a record of where and when liquid water was capable of transporting sediment on Mars in the past. Furthermore, as depositional features in close proximity with their sediment sources, alluvial fans preserve a record of environmental changes that occurred while sediment was being transported and deposited. In this study, we leveraged high resolution images from the CTX camera on Mars Reconnaissance Orbiter, which cover the entire surface of Mars, to conduct a global survey for fan-shaped sedimentary landforms across the martian surface and explore their implications for environmental change.
Image 1: A large alluvial fan in Harris crater, centered on 21.5°S, 67°E. The down-fan-trending ridges are interpreted to be paleochannels now in inverted relief due to preferential erosion of finer overbank deposits. This particular fan exhibits evidence for multiple episodes of sedimentation (Williams et al., 2011). CTX image G02_019160_1580_XN_22S292W.
Post contributed by Dr. Sierra Ferguson. Postdoctoral Researcher, Department of Space Studies, Southwest Research Institute, Boulder Colorado, USA.
Dione is one of Saturn’s many photogenic “mid-sized” icy moons, first visited by the Voyager mission in 1980. Most well-known for its system of bright scarps/troughs on the trailing hemisphere of the moon (aka the “wispy terrain”, Image 1), Dione is host to a wide range of interesting surface and subsurface features such as impact craters, tectonics, and potentially a subsurface ocean. Models of the formation of the Saturnian satellites have frequently placed their formation ages at ~ 4 billion years ago, roughly around the same time that Saturn itself formed. However, recent modeling of the orbital dynamics of the system have shown that the inner moons (Mimas, Enceladus, Tethys, Dione, and Rhea) may in fact be much younger, with the youngest potential formation age only 100 million years ago. One way to examine the surface ages, and the potential formation age of Dione, is through the analysis of impact crater distributions. We utilized images from NASA’s Cassini spacecraft to examine the sources of craters on Dione and what they mean for the ages of Saturn’s satellites and the evolution of Dione’s surface (Ferguson et al., 2022b).
Impact craters are very common geologic feature on planetary surfaces. We can analyze their distributions to examine their origins, and model the surface age of a mapped region. A lack of impact craters is often interpreted to be a result of crater erasure by geologic resurfacing process (i.e., volcanism, mass wasting, additional cratering). By mapping the observed craters on Dione, we are able to characterize the bombardment environment at Saturn as well as the modification history of Dione’s surface.
Image 1: Cassini ISS-NAC image of Dione’s wispy terrain area. This image provides context for the craters on the surface as well as the relationship of the craters to the tectonic features observed on the surface. Image credit NASA/JPL/SSI, PIA 18327 (https://photojournal.jpl.nasa.gov/catalog/PIA18327)
Post contributed by Dr. Costanza Rossi, Astronomical Observatory of Padova, Italian National Institute for Astrophysics
During their 1979 flybys of the Jovian system, the Voyager probes revealed to us abundant evidence of global-scale structures on the icy satellite Ganymede (Image 1), showing that tectonic activity played a key role during its geologic history. Ganymede’s puzzling surface covers an internal liquid ocean and is characterized by two main geological units, the dark and light terrains (> 4 billion and 2 billion years old, respectively; Patterson et al., 2010), which are shaped by a complex tangle of linear to curvilinear morpho-tectonic structures (Image 1). There are brittle structures called furrows within the dark terrain, and grooves within the light one, which (apparently) represent disconnected deformation histories. Impact cratering processes affected the dark terrain by forming furrows with graben-like morphologies, which are interpreted as remnants of a multi-ring basin formed by a giant ancient impact (e.g., Prockter et al., 2000). On the other hand, the light terrain formed at the expense of the dark one, as a result of strong tectonics that formed the grooves (e.g., Pappalardo et al., 1998). Although extension is considered as the main process for the groove formation, it is assumed that strike-slip has been also pivotal during their development (Cameron et al., 2018; Rossi et al., 2018). The lack of pure compressional structures still represents one of the open enigmas of the tectonics of Ganymede. Ganymede’s tectonic-related landforms potentially represent pathways connecting the surface with the underlying ocean, and their investigation is pivotal to understand their possible development and interconnection at depth (e.g., Lucchetti et al., 2021), their evolution, and the endogenic processes responsible for their formation.
Image 1: Orthographic projection of the leading hemisphere of Ganymede (Voyager/Galileo global mosaic of Kersten et al., 2021). The dark squares show the location of Image 2 and the white square shows the location of Image 3.
Post contributed by Dr. Petr Brož, Institute of Geophysics of the Czech Academy of Sciences.
Kilometer-sized flows (KSFs) have been observed in many regions on Mars and have typically been interpreted as lava flows. However, sedimentary volcanism has been proposed as an alternative origin for some martian KSFs. Remarkable examples of such hypothesized sedimentary KSFs are located at the southern margin of Chryse Planitia, a large impact basin that has been partly infilled by sedimentary strata. There, dozens of such flows (Images 1 and 2) are associated with >1000 conical and dome-shaped edifices that have been previously interpreted to result from subsurface sediment mobilization (Komatsu et al., 2016, Brož et al., 2019); however their formation mechanism still remains enigmatic due to the absence of ground truth. To get deeper insight about the formation of these KSFs we generated seven Digital Elevation Models computed from HiRISE stereo pair images enabling us to reveal their detailed, meter-scale morphology (Image 2). As a result, we have been able confirm that KSFs are indeed aggradational edifices formed by the movement of a low viscosity liquid (Image 2b). Several of the observed morphologies associated with these flows suggest that such liquid had to be a mixture of water and sediment, i.e., mud, rather than silicic lava. In such a case, the southern part of Chryse Planitia is a region on Mars where subsurface sediment mobilization could have operated in the past and hence represents a promising site for future exploration where deeper-sourced sedimentary deposits are exposed at the surface.
Image 1: An example of kilometer-sized flow in Chryse Planitia consisting of three main parts as captured in CTX image P17_007639_1997_XN_19N034W, centered 19.85°N, 326.02°E).
Post contributed by Emily Bamber (PhD Candidate), Jackson School of Geoscience, University of Texas at Austin
Impact craters are a common feature of the surface of Mars and other planetary bodies. Impact craters are formed by the force of an asteroid or comet colliding with the surface of a planet. This collision excavates a bowl-shaped depression and causes uplift, deformation and deposition of excavated material to form a crater rim, which can rise hundreds of meters above the surroundings (for a crater on the scale of tens of kilometers in diameter). However, some of Mars’ impact craters stand out from their counterparts on other rocky planets and moons because they have been affected by the action of erosion by liquid water. In comparison to the Moon, Mars’ craters deviate much more from their original geometries, but perhaps more strikingly, more than 400 craters on Mars’ surface have an inlet valley (Image 1); a valley which starts from outside the crater and crosses the rim crest (e.g. Cabrol & Grin, 1999; Goudge et al., 2016). Although no substantial liquid water currently flows on Mars today, the preserved valleys likely delivered water from a crater’s surroundings to the crater interior, at some time in the distant past. Therefore, their formation could give clues to Mars’ watery past.
Image 1: An impact crater on Mars, at 16.2°N, 53.2°W with a high-standing crater rim around its circumference. To the south of the crater, there is an inlet valley that crosses the rim divide, with two wide tributaries. The valleys are now dry, but liquid water in Mars’ past likely incised them. The crater does not have an outlet valley (where water may have exited the crater in a through-flowing system). Background is a CTX DEM (from the stereo pair B21_017806_1955 – B21_017951_1955), HRSC DEM h2277_0000 and global MOLA DEM, overlain on CTX image B21_017806_1955 and THEMIS global IR mosaic.
Post contributed by Dr. Edgard G. Rivera-Valentín, Lunar and Planetary Institute, Universities Space Research Association.
Although its surface can reach 800°F (427°C), some of Mercury’s craters conceal vast ice deposits that, in a sense, “sparkle” in the light of radar (Image 1). The so-called radar bright features were first identified in 1992 using ground-based observatories, in particular the Arecibo Observatory in Puerto Rico, which provided magnificent views down to craters tens-of-kilometers in diameter all the way from Earth. The ice lies within craters whose morphology and location results in areas that do not receive direct sunlight (i.e., permanently shadowed regions). This allows for the low temperatures needed to retain ice over millions of years.
Image 1: Radar image of Mercury’s north polar terrain (> 75°N) in polar stereographic projection. The five notable craters, Chesterton (88.5°N, 126.9°W), Tolkien (88.8°N, 211.1°W), Tryggvadóttir (89.6°N, 171.6°W), Kandinsky (87.9°N, 281.2°W), and Prokofiev (85.7°N, 297.1°W) are labeled. Radar backscatter is noted in grayscale from black (below noise levels) to white (high backscatter). All bright areas are radar bright features, which are located within craters that have permanently shadowed regions. These are the locations of ice deposits on Mercury. For scaling reference, Chesterton crater has a diameter of 37.2 km. Image credit: Figure 2 in Rivera-Valentín et al. (2022) PSJ 3,62.
The discovery of the radar bright features at Mercury’s poles was one of the motivators of NASA’s MESSENGER (Mercury Surface, Space Environment, Geochemistry, and Ranging) mission, which studied Mercury between 2011 – 2015. That is because although the radar scattering properties of the features were reminiscent of observations of the Martian polar layered ice deposits, as well as of the icy moons of Jupiter, they alone did not uniquely indicate the presence of water ice. The detailed studies by MESSENGER along with the earlier radar observations together now strongly suggest deposits of water ice. This is due to their location, evidence from high resolution and long exposure imaging, and measurements of epithermal and fast neutron fluxes.
Post contributed by Dr. Matthew Chojnacki, Planetary Science Institute, Lakewood, CO, USA.
Megaripples, which are intermediate-scale aeolian bedforms as compared with ripples and dunes, have been studied extensively on Mars and thought to be largely inactive landforms from past climates. On Earth, “megaripples” can be 30 cm to tens of meters in spacing with an abundant coarse sand population resulting in reduced migration except for very strong storm events. We mapped the extensive sand seas across the north pole of Mars for the presence of Martian megaripples (Image 1).
Image 1. HiRISE perspective view of north polar dune fields with active megaripples (arrows). Megaripple activity is most evident on the upwind edges of dune fields and in some cases within inter-erg areas. Clusters of contiguous megaripple fields often flanked by ripples and dunes were most common, while occasionally occurrences of mobile megaripple trains atop of bedrock were observed. All images and associated digital terrain models can be accessed at: https://hirise.lpl.arizona.edu/ Credit: NASA/JPL/University of Arizona.
Post contributed by Dr Joel Davis, Department of Earth Sciences, Natural History Museum, UK
On Earth, the interplay between tectonics and sedimentation is commonplace. In regions such as Death Valley, USA, extensional tectonic processes have formed sedimentary basins, where the subsequent erosion of the landscape causes the formation of features such as alluvial fans. Alluvial fans are conical-shaped, sedimentary deposits usually formed by stream flows in upland regions. Later tectonic activity can lead to the deformation and relative uplift of these landforms. A new generation of alluvial fans may then form, overprinting the older generation and extending down to the new position of the basin floor.
To contrast, as most alluvial fans on Mars are found within impact craters, they do not show evidence of tectonic deformation. The alluvial fans hosted within the canyons of Valles Marineris appear to be an exception to this (Figure 1). Valles Marineris may have developed over billions of years by repeated episodes of normal faulting. In several regions of Valles Marineris, not only are alluvial fans found on the current floor, but remnants of such features are found on the canyon walls as well – perched several kilometres above the floor!
Figure 1: Faulted alluvial fan deposits in Coprates Chasma, Mars. (A) Oblique Context Camera (CTX) image of multiple generations of alluvial fan deposits, offset by faulting. Credit: NASA/JPL/MSSS. Inset shows High Resolution Imaging Science Experiment (HiRISE) image of bright-toned layering in fan margins. Topographic contour interval is 200 m. Credit: NASA/JPL/UoA. (B) HiRISE mosaic of C2 alluvial fan deposits on canyon floor, where surfaces contain distributary channels and paired ridges, cut by multiple faults. Youngest C3 fans superpose faults and are local to C2 fan apices. Credit: NASA/JPL/UoA. (C) Planform CTX mosaic of oldest (C1) alluvial fan deposits (perched terraces), which have been faulted and later eroded (orange line). NASA/JPL/MSSS. After Davis et al., 2021.
Cones are common landforms arising from volcanism on Earth and can be observed on other terrestrial planets. Mars hosts small conical hills known as mounds – these positive relief features have been reported in several locations over the Martian surface (Brož et al., 2017; Hemmi & Miyamoto, 2017) and inside craters such as Firsoff, Becquerel, Kotido, and Crommelin in southwestern Arabia Terra (cf. McNeil et al., 2021; Pondrelli et al. 2015). Various processes have been proposed for the development of mounds on Mars, including mud volcanism, monogenetic or small-scale volcanism, preferential erosion of sedimentary units, and diapir-like discharge of sediment-laden fluid. Sedimentary origins have been previously hypothesized for the mounds found on crater floors (Pondrelli et al., 2011; Pondrelli et al. 2015). In that context, we identify previously unrecognized evidence of intrusive igneous process at the center of an unnamed ∼85-km diameter floor-fractured crater (FFC, Bamberg et al., 2014) in North-Central Arabia Terra (Image 1). We consider whether the observed geomorphic features are related to impact-induced volcanism or to regional intrusive activity during the late Noachian to early Hesperian periods of Mars (~3.6 Ga; Rani et al., 2021).
Image 1:Thermal Emission Imaging System daytime Infrared image mosaic of the studied Floor-Fractured Crater in the eastern Arabia Terra region, overlain by color-coded Mars Orbiter Laser Altimeter and High-Resolution Stereo Camera elevation map. The black box in the center of the crater shows the location of the mounds and ridges (See Image 2). The location of the crater is marked by a yellow star in the inset global map of Mars. Image credits: HRSC-MOLA Blended DEM Global 200m v2, USGS/DLR/NASA/Goddard Space Flight Center. THEMIS-IR Day Global Mosaic 100m v12, THEMIS team/ASU
Post contributed by Meven Philippe, Ph.D student, Laboratoire de Planétologie et Géodynamique de Nantes, CNRS/Université de Nantes, France.
On Earth, sharp drops of temperature below 0°C can cause the ground to contract and fracture, forming thermal-contraction cracks. These cracks can connect in a network, which leads to a ground that shows a polygonal pattern. Ice or sand can then seasonally aggrade as wedges into these cracks, which uplifts polygons’ margins to higher elevations than their centres, thus forming low-centered polygons. When the infilling material degrades, the elevation of the margins reduces with respect to that of the centre, thus forming high-centered polygons. The same processes are assumed to occur on Mars in the mid-latitudes. Image 1 (42.9°N, 115.43°E) is located in Utopia Planitia, a lowlands plains area in the martian Northern hemisphere where numerous potential periglacial morphologies have been described.
Image 1: Polygonal patterned ground in a crater in Utopia Planitia, Mars. The blue circle highlights examples of low-centred polygons, the red circle highlights examples of high-centred polygons. The black arrow points out the sharp border between the polygonised unit and another one, non-polygonised. Illumination is from the left. Credits: HiRISE image ESP_036366_2235 – NASA/JPL/University of Arizona.
Post contributed by Prof. Paul K. Byrne, Department of Earth and Planetary Science, Washington University in St. Louis, USA.
Venus is a tectonically tortured world (Image 1). Vast rift systems, widespread crustal shortening, and a crumpled terrain type, tessera, collectively attest to major tectonic activity on a world only a little smaller than Earth. In places, strains are spatially distributed; in others, extension or shortening is concentrated into narrow bands. And, where these bands intersect, they define a type of tectonic deformation on Venus that hasn’t been recognised before.
Image 1: A 1,100 km-wide, false-colour radar view of Lavinia Planitia, one of the lowland regions on Venus where the lithosphere has fragmented into blocks (purple) delineated by belts of tectonic structures (yellow). Image credit: Paul K. Byrne.
Post contributed by Adomas Valantinas, Physikalisches Institut, Universität Bern, Switzerland.
Slope streaks are linear, fan-shaped and sometimes anastomosing dark markings that form frequently on equatorial Martian slopes. Image 1 shows examples of these features imaged by the Colour and Stereo Surface Imaging System (CaSSIS) on board ExoMars Trace Gas Orbiter. Slope streaks are found on steep slopes but can also initiate on gentle slopes well below the angle of repose. They can propagate for many hundreds of meters and pass around obstacles. Unlike many other seasonal features on Mars, which form and fade periodically, slope streak markings remain visible for several decades after their formation. Some dark streaks have been observed to transition into bright streaks over time. Therefore, these features represent a unique phenomenon that provide important insights into Mars’ dynamic surficial and atmospheric processes.
Image 1: Bifurcating morphology of dark slope streaks on a crater wall in the Mars’ equatorial latitudes (9.19°N, 39.60°E). Note the streaks’ ability to go around obstacles, their sharp boundaries, and their digitate termini. CaSSIS ID: MY36_015777_009_1. Credit: ESA/Roscosmos/University of Bern.
Post contributed by Dr Colman Gallagher, UCD School of Geography and UCD Earth Institute, University College Dublin, Ireland.
Viscous flow features occur in the mid-latitudes of Mars and have characteristics consistent with being debris-covered glaciers. Climate models suggest that martian glaciers are cold-based systems in which meltwater has never been widely produced but Gallagher et al. (2021) described a widespread assemblage of landforms in Phlegra Montes (Image 1) that are morphologically analogous to landforms classified by Glasser and Bennett (2004) as indicative of warm-based glacial erosion on Earth. This suggests that glaciers in this region produced meltwater in the past. Warm-based glaciers reach the melting point at their beds, allowing them to slide over their beds and erode the underlying landscape. Meltwater produced by warm-based glaciers can also flow across the bed and erode subglacial meltwater valleys. Landforms of glacial erosion identified by Gallagher et al. (2021) include zones of aerial scour, glacial troughs (Image 2), corries (cirques; Image 3), lee-side rock cavities, streamlined bedrock forms, rock grooves (Image 4), subglacial meltwater channels (Image 4), and tunnel valleys. This is in addition to a depositional glacial meltwater landform called an esker, which was discovered by Gallagher and Balme (2015) and you can read about in a previous post here.
Image 1. Elevation map of Mars showing the location of Phlegra Montes. Blue regions are low elevation and red/white regions are high elevation (MOLA elevation data credit: USGS, NASA, JPL).
Post contributed by Dr. Alexander Barrett, School of Physical Sciences, The Open University, UK.
Transverse Aeolian Ridges (TARs) are a distinctive aeolian landform, commonly found on the surface of Mars (Balme et al., 2008; Bourke et al., 2006). They consist of metre to decimetre scale granular ripples, which form transverse (i.e. perpendicular) to the direction of the prevailing wind. Similar to terrestrial megaripples, TARs are believed to form through surface creep of coarse particles. Understanding the distribution of TARs is important for understanding the aeolian environment in which they formed (Favaro et al., 2021), and for planning rover missions (Balme et al., 2018). TARs present a potential hazard when they occur in rover landing sites. Consequently, a team at the Open University has trained a machine learning model called NOAH-H (developed by SciSys Ltd. for the European Space Agency; Woods et al. 2020) to identify these features, and distinguish them from other common surface textures (Barrett et al. 2021).
Image 1 shows bedforms in Oxia Planum. The size and prevalence of bedforms in this area is important, since it is the landing site for the ExoMars Rosalind Franklin Rover (Vago et al., 2017), which will arrive in 2023 to search for signs of past and present life in this once water-rich area (Quantin-Nataf et al., 2021). Large bedforms, or continuous regions of loose aeolian sand, could form navigational hazards, limiting the directions in which the rover can progress, or the science targets that it can access. We need to have a good understanding of the distribution of TARs in the immediate area where the rover will land, however this is not precisely known in advance and could be anywhere within a large landing ellipse. Mapping all of the TARs in the potential landing area by hand is impossible in the time available, even for a large team.
Image 1: Large TARs at the planned ExoMars Rosalind Franklin rover landing site at Oxia Planum, Mars (High Resolution Imaging Science Experiment (HiRISE) Image ESP_048292-1985. a) translucent layer showing Machine Learning produced terrain classification, overlying the red-band HiRISE image from which the classification was made. The area is dominated by rugged and fractured bedrock, while topographic lows contain non-bedrock material. This has been shaped into TARs of various sizes and degrees of continuity. The divide between bedrock and non-bedrock textures is well defined, as are the perimeters of patches of TARS. The transitions between textured, fractured and rugged bedrock are less well identified, as these features form a continuum, and so some areas have characteristics of multiple classes. HiRISE Image Credit NASA/JPL/University of Arizona.
Post contributed by Charlene Detelich, Department of Geological Sciences, University of Alaska Anchorage, USA.
Voyager flybys of the Saturnian system in the early 1980s discovered a ring of mountains on Saturn’s icy, two-toned moon, Iapetus. These mountains are up to 20 km tall, over twice the height of Mt. Everest, yet Iapetus is almost 650 times smaller than Earth. Later flybys from the Cassini spacecraft (Image 1) revealed these mountains to be part of a ridge encircling nearly all of Iapetus’s equator, resulting in the moon’s bizarre walnut shape. Numerous endogenic (internal) processes have been proposed for the origin of the equatorial ridge, including tidal spindown, convection, despinning coupled with global volume change, global contraction, or intrusion. However, a recent study by Detelich et al. (2021) found that evidence from photogeological mapping and crater statistics support an exogenic (external) origin best explained by the accretion of an ancient ring system onto Iapetus’s surface.
Image 1: Cassini image PIA08376 (https://photojournal.jpl.nasa.gov/catalog/PIA08376) taken on September 10th, 2007 shows the enigmatic equatorial ridge of Saturnian moon, Iapetus. The ridge runs vertically down the center of the image before disappearing into the shadows of the non-lit portion of Iapetus. Image enhanced by Charlene Detelich.
Post contributed by Dr. Elodie Lesage, Laboratory of Geosciences Paris-Saclay, University of Paris-Saclay.
The young (60-100 Myrs) and active surface of the Jupiter’s icy moon Europa includes various geological features. In some locations, smooth plains have been observed on Europa, defined as features with no visible texture and an albedo lower than the surrounding terrains. The smooth plains overprint the preexisting terrains, and are confined within basins bounded by topographically high features.
Image 1: Four smooth plains on Europa. The arrows show the sunlight direction. Galileo images: (a) Image 5452r, resolution: 27 m/px; (b) Image 0713r, resolution: 57 m/px; (c) Image 0739r, resolution: 57 m/px; (d) Image 9352r, resolution: 60 m/px.
The best-resolved images of Europa’s surface were acquired by the Galileo spacecraft between 1996 and 2001. Image 1 shows four smooth plains visible in Galileo images, which share the following common characteristics: i) thin features occupying topographic lows, ii) a smooth appearance with little or no visible texture, and iii) kilometer-scale width with a quasi-circular, lobate shape. The morphology of these features suggests that they result from the flow of low-viscosity fluid, such as liquid cryomagma (i.e., briny water coming from Europa’s interior).
Post contributed by Dr. Lauren Mc Keown, School of Physical Sciences, The Open University, UK.
Spiders are unusual branched landforms found among the high southern latitudes of Mars (Image 1). They have no Earth analogues and are often accompanied by fans and spots that appear in spring. They are proposed to form when sunlight penetrates the Martian south polar seasonal CO2 ice layer, causing ice at its base to change from ice to gas, and eventually crack. Escaping gas then scours the terrain beneath, carving spider-like patterns and depositing material on top of the ice via a plume (Kieffer et al., 2003). However, although this suggested process is well-accepted, it has never been directly observed on Mars. In order to investigate whether spider patterns could form by CO2 sublimation under Martian atmospheric pressure, experiments were performed (Image 2) at the Open University Mars Simulation Chamber, which simulates Martian atmospheric conditions.
Image 1: Examples of spiders on Mars (HiRISE image ESP_014282_0930). Left shows the ‘classic’ spider morphology which consists of a central depression and radial tortuous dendritic troughs emanating from its centre. Right is a context image. Image Credit: NASA/JPL/University of Arizona.
Post contributed by Dr. Alfred S. McEwen, Lunar and Planetary Laboratory, University of Arizona, USA.
Recurring slope lineae (RSL) are dark linear markings on steep slopes of Mars that regrow annually and likely originate from the flow of either liquid water or dry granular material. Following the great dust storm (or planet-encircling dust event) of Mars Year 34 (in 2018), the High Resolution Imaging Science Experiment (HiRISE; McEwen et al., 2007) on Mars Reconnaissance Orbiter (MRO) has seen many more candidate RSL than in typical Mars years (Image 1). These RSL sites show evidence for recent dust deposition and dust devil activity, so dust lifting processes may initiate and sustain RSL activity on steep slopes.
Image 1: RSL and dust devil tracks on a hill in the southern middle latitudes (41.1ºS, 187.4ºE). Inset shows some of the RSL at higher resolution. Hundreds of dark dust devil tracks are seen on the full image as the more diffuse lines that cut across topography. HiRISE image ESP_058122_1385, acquired after the 2018 dust storm. Credit: NASA/JPL/University of Arizona
Post contributed by C. Adeene Denton, Department of Earth, Atmospheric, and Planetary Sciences, Purdue University, USA.
Antipodal terrains are unusual regions of hilly, lineated, or otherwise disrupted terrain that are on the direct opposite side of planetary bodies to large impact basins. These mysterious terrains have been observed at the antipodes to the Caloris basin on Mercury and the Imbrium basin on the Moon, where their formation is considered to be indicative both of the impact’s size and the specificities of the planetary body’s interior structure. Recent revisiting of data from the New Horizons spacecraft revealed an unusual region of disrupted and lineated terrain on Pluto’s far side that is roughly antipodal to the massive Sputnik Planitia basin, the feature sometimes referred to as “Pluto’s Heart” (Image 1). If the lineated terrain is indeed connected to the large impact believed to have formed Sputnik Planitia, then the two geologic features offer a new and unusual way to probe Pluto’s interior: seismology through giant impact.
Image 1: Comparison of Pluto’s nearside (left) and farside (right) with Sputnik Planitia and its proposed antipodal terrain indicated. The location of Image 3 is also indicated. Images modified from full-scale planetary images taken by the New Horizons spacecraft, via NASA/JHUAPL/SWRI.
Post contributed by Divya M. Persaud, University College London, UK/NASA Jet Propulsion Laboratory, USA.
This canyon, Sakarya Vallis, cross-cuts the central mound of Gale Crater, Mars, and was probably formed by fluvial erosion. Gale Crater, the exploration site of the Curiosity rover, has undergone a complex geologic history of aqueous and aeolian processes. The central mound is a topographic high in the center of the crater, on which the ~5.2 km peak Aeolis Mons is situated. This feature sports several canyons (which cut through it), yardangs, and spectacular exposed layers, and its origins are uncertain (likely aeolian and/or fluvio-lacustrine). Image 1 shows the largest canyon on the mound, at 26 km long, up to 3.5 km wide, and up to 400 m deep. The hundreds of meters of exposed layers in this canyon provide a glimpse into the depositional history of the central mound of Gale crater.
Image 1: A 3D perspective view of the interior of the channel, visualized from HiRISE Digital Terrain Model DTEEC_006855_1750_007501_1750_A01 and HiRISE image PSP_007501_1750. This view shows the exposed layers and a possible fracture on the eastern wall (left), while the topography of the bright, inverted feature can be seen on the floor of the canyon. The floor of the canyon is overlaid by ripples. The scalebar is oriented north.
The center of Sakarya Vallis is approximately 27 km from and 700 m higher in elevation than the base of Gediz Vallis, the approximate location of the Curiosity rover (as of February 2021, sol 3038). The layers exposed in Sakarya Vallis therefore represent later depositional events than those encountered by Curiosity to date. As the rover ascends the mound towards Aeolis Mons, this geology could help contextualize rover observations and constrain lateral differences in palaeoenvironment.
The canyon cross-cuts layered hydrated sulfates in the lower mound, while spectra of the upper mound point to a dust composition. The surface of the upper unit (Image 2) is an etched, yardang unit. Yardangs are streamlined features eroded by the wind.
Image 2: A) An overview of the canyon on the central mound from CTX imagery. The slope of the canyon is to the northwest. The surface of the mound is etched by yardangs and in-filled craters. B) A close-up view of the northern, lower half of the canyon, from HiRISE image PSP_007501_1750. The extent of the inverted feature can be seen on the floor.
Within the canyon are possible point bars formed by past rivers (Image 2), mass-wasting features (Image 3), ripples (Image 1), and boulder-scale lag deposits along the floor and clifftops. A bright, topographic feature resembling an inverted channel or meander belt (Image 1) extends along much of the canyon floor and may represent later flow of liquid water, and subsequent topographic inversion by erosion of materials surrounding the channel.
Image 3: Layers exposed in the northern part of the canyon in a possible mass-wasting feature (HiRISE image PSP_007501_1750). These layers are sub-horizontal and dip gently to the northwest. The ripples on the floor have wavelengths of 10-20 m.
Post contributed by Axel Noblet, Laboratoire de Planétologie et Géodynamique de Nantes, CNRS/Université de Nantes, France
Jezero Crater on Mars will soon be explored by NASA’s Perseverance rover. This crater has been interpreted as a paleolake. It contains two fan-shaped deposits in the northern and western portions of the crater. These deposits have been interpreted as ancient deltas. The delta located in the western portion of Jezero (Image 1) displays some of Mars’ best-preserved fluvio-deltaic features, and exhibits a variety of structures such as inverted channels and point-bar strata (Image 2). This delta contains a precious record of various depositional environments, and in situ exploration can give us insight of Mars’ fluvio-lacustrine history. The association of well-preserved lacustrine features with orbital detections of carbonates along the inner margin of Jezero points toward high biosignature preservation potential for these deposits. Hence Jezero’s western delta contains a record of the evolution of Mars’ ancient climate and possible habitability. The presence of a long-lived lake system on Mars is astrobiologically significant, and the deposits within the Perseverance landing site could have preserved biosignatures that could be investigated and cached for a future sample return mission.
Image 1: 3D view of Jezero western delta, looking north from the center of the crater. The data visualized here is a CTX camera orthorectified mosaic draped over a CTX digital terrain model (horizontal resolution: 20m/px). The triangular ‘birdfoot’ shape of the delta is clearly visible, and inverted channels can be seen radiating from the apex of the delta. The inlet valley goes diagonally from the upper left of the image through the delta deposits. The crater floor appears as the smooth terrains on the lower part of the image.
Post contributed by Dr. Caitlin Ahrens, NASA Goddard Space Flight Center, USA.
Cryovolcanism involves the transfer of icy or gaseous subsurface materials either to the surface (eruptive) or through the subsurface (non-eruptive) of an icy planetary body. It differs from magmatism and volcanism on Earth, which involves the migration and eruption of molten rock. Cryovolcanism is thought to have operated on several icy bodies in the Solar System, including Enceladus, Triton, Pluto, and possibly Europa. Cryovolcanism results primarily from internal heat-producing processes, and excludes sublimation and condensation processes at the surface. In the case of Pluto, there is evidence for a subsurface fluid layer, the presence of cryovolcanoes, and cryovolcanic subsurface materials (called cryomagma) which can contain ammonia and methane. Due to the presence of a deformable subsurface layer, it is possible for the material to shift, causing uplift followed by a collapse-type event. This is a possible scenario at Hekla Cavus (Image 1), a large, elongated, and irregular depression situated within a much larger north-south (N-S) ridge-trough system outlined by mountain ranges.
Image 1: Image of Hekla Cavus taken from the LORRI instrument onboard the New Horizons.
Post contributed by Dr Jack Wright, School of Physical Sciences, The Open University, UK.
The Caloris basin is the largest (~1,500 km across), well-preserved impact structure on Mercury (Image 1a; Fassett et al., 2009). Hummocky plains around Caloris host numerous, steep-looking, conical knobs (Image 1b). The obvious explanation for the hummocky plains is that they formed from material ejected by the Caloris impact ~3.8 billion years ago. It follows that the knobs probably formed from discrete ejecta blocks. What isn’t obvious is why many of these blocks, which hypothetically could have formed with a variety of shapes, exist as steep cones in the present day. If these knobs really did form as Caloris ejecta, then they offer a rare opportunity to study materials ejected from Mercury’s interior with remote sensing techniques.
Image 1: Mercury and the circum-Caloris knobs. (a) Enhanced colour limb view of Mercury from the MESSENGER spacecraft. The Caloris basin’s interior is made of volcanic plains that appear orange in this data product. The arrow indicates the location of (b). (b) Examples of circum-Caloris knobs just outside the Caloris rim. Mosaic of MESSENGER MDIS WAC frames EW0220807059G, EW0220807071G, and EW0220763870G. ~86 m/pixel.
Gullies are found on steep slopes on the surface of Mars and appear as a linear-to-sinuous channel linking an alcove at the top to a fan at the bottom. The most interesting interpretation of the past two decades has been that the Martian gullies were carved by the flow of liquid water as discovered from the high-resolution images returned by the Mars Orbiter Camera (MOC) onboard “Mars Global Surveyor (MGS)” in 2000 (Malin and Edgett, 2000). Subsequent observations using MOC and the Mars Reconnaissance Orbiter (MRO) High Resolution Imaging Science Experiment (HiRISE) images revealed that Martian gullies are active today and that sublimation of seasonal carbon dioxide frost – not liquid water – could have played an important role in their formation. In our recent work using HiRISE images we reported global distribution of overlapping lobate deposits in gullies (Image 1) showing that a debris-flow like process may be responsible for gully formation (Sinha et al., 2020).
Image 1: Top: 3D view of gullies on the pole-facing wall of ~8 km diameter Los crater (35.08˚ S, 76.22˚ W) on Mars. HiRISE image ESP_020774_1445 (0.25 m/pixel) is draped over 1 m/pix HiRISE elevation model. The depth of crater floor from the crater rim is ~1 km in elevation and the image spans ~4 km from left to right. The box shows the location of bottom panel. Bottom: Image shows the gully fan surface within Los crater with overlapping lobate deposits, including convex-up and tongue shaped terminal lobes with lateral levees. HiRISE image credits: NASA/JPL/University of Arizona.
Some impact craters are classified as polygonal impact craters (PICs), which have at least one straight rim segment, as shown in Image 1. The morphologies of PICs are shaped by pre-existing, sub-vertical structures in the target material, such as normal and strike-slip faults, joint sets, and lithologic boundaries. Because the straight rim segments of PICs only form where pre-existing structures are present, PIC morphologies can be used to analyze fractures that are buried by regolith or too small to be seen in available spacecraft images. On the icy Uranian moon Miranda, PICs are widespread across its southern hemisphere, which was imaged by the Imaging Science System (ISS) onboard the Voyager 2 spacecraft. Some of these PICs reveal previously undetected fractures that suggest Miranda has experienced multiple periods of tectonic activity.
Image 1: Examples of two PICs identified on the Uranian moon Miranda. Black arrows indicate the straight rims of these PICs. The Voyager 2 ISS image mosaic shown here includes the following images, from top to bottom: c2684620 (light blue box), c2684629, c2684617 (dark blue box).
Images of the lunar surface reveal layered deposits presumed to be sequences of basaltic lava flows. These sequences have been imaged since the Apollo astronauts acquired both orbital and surface photographs in the 1960s and 1970s. Apollo 15 astronauts visited Hadley Rille, a 130 km long, 200 m deep sinuous feature that was formed by flowing lava, similar to lava channels or tubes on Earth. Photographs taken by the astronauts (such as Image 1) show that the rille cut into the underlying substrate, revealing sequences of layered material. The layers are believed to be basaltic lava flows, based on outcrop morphologies and nearby samples. The thicknesses of ancient lava flows provide insight into the emplacement, dynamics, and history of volcanism on the Moon.
Image 1: Apollo 15 surface image of the interior wall of Hadley Rille (https://www.hq.nasa.gov/alsj/a15/AS15-89-12106HR.jpg). Inset highlights layered deposits presumed to be basaltic lava flows with possible intercalated regolith deposits. Outcrop is approximately 8 meters thick.
Post contributed by Helle L. Skjetne, PhD candidate, Department of Earth and Planetary Sciences, University of Tennessee, Knoxville, USA.
Chaos terrain is formed by disruption of preexisting surfaces into irregularly shaped blocks with a “chaotic” appearance (Image 1). This typically occurs through fracturing (that can be induced by a variety of mechanisms), and the subsequent evolution of these blocks can follow several paths (Image 2). These distinctive areas of broken terrains are most notably found on Jupiter’s moon Europa, Mars, and Pluto. Although chaos terrains on these bodies share some common characteristics, there are also distinct morphological differences between them (Image 1). The geologic evolution required to shape this enigmatic terrain type has not yet been fully constrained, although several chaos formation models have been proposed. We studied chaotic terrain blocks on Pluto, Europa and Mars to infer information about crustal lithology and surface layer thickness (Skjetne et al. 2020).
Image 1: Examples of chaotic terrain “blocks” (referring to each mountain-like topographic feature). Chaos on Pluto in a) Tenzing Montes and b) Al-Idrisi Montes, respectively (New Horizons image at ~315 m px–1), c) Hydraotes Chaos on Mars (Mars Odyssey THEMIS daytime infrared global mosaic at 100 m px–1), and d) Conamara Chaos on Europa (Galileo 210–220 m px–1 East and West RegMaps).
Post contributed by Adam J. Hepburn, Department of Geography and Earth Sciences, Aberystwyth University, UK.
Mars hosts abundant glacier-like landforms throughout its mid-latitudes, the presence of which necessitates major shifts in climate relative to present conditions. These ice-rich viscous flow features (VFFs) are typically found in coalescing, size-hierarchical systems whereby lower-order glacier-like forms (GLFs; ~5 km long) flow from alcoves and merge with higher-order lineated valley fill (LVF; 100s of km long). Several larger VFFs have been dated previously, indicating Mars underwent glaciation in the past several hundred million years, during the late Amazonian epoch. However, several authors have noted examples of GLFs flowing onto, rather than into, LVFs (Image 1), and hypothesised that these may correspond to a more recent phase of glacial activity. We used crater dating to ascertain that—in addition to the earlier phase of widespread regional glaciation—these superposed GLFs (SGLFs) were formed following at least two major cycles of more recent alpine glaciation, the latter of which ended ~2 million years ago.
Image 1: Superposed glacier-like form (SGLF) flowing onto the underlying viscous flow feature (underlying VFF), in the Protonilus Mensae region of Mars. (A–B) North-up orientated HiRISE image (ESP_018857_2225) image of an SGLF (light blue) emerging from an alcove and flowing onto lineated valley fill (dark blue). Approximate location of image centre is 42.23◦ N, 50.53◦ E. Reproduced from Hepburn et al, 2020.
Post contributed by Michael J. Malaska, PhD, Scientist, Jet Propulsion Laboratory / California Institute of Technology, USA.
Saturn’s moon Titan is where organic chemistry and surface geomorphology intersect to create an enigmatic landscape with many features in common with Earth, but that are made of completely different materials. Much of Titan’s surface is made up of organic sedimentary materials; recent mapping shows that plains and dunes cover over 80 percent of the globe. The Cassini spacecraft’s Synthetic Aperture Radar (SAR) was able to penetrate Titan’s thick haze and reveal areas of highly dissected plateaux on the surface that are called labyrinth terrain. Image 1 shows an SAR image of an example of this type of terrain, the Sikun Labyrinth. Detailed examination of Titan’s labyrinth terrain can tell us a lot about Titan’s geological history and surface evolution.
Image 1. Top: Image of the Sikun Labyrinth in the south polar terrain of Titan. The blue arrow and number at top left indicates direction of radar illumination and incidence angle for this scene. Bottom: diagram showing how radar illumination interacts with terrain of valleys and plateaux. Image credit: Mike Malaska.Read the full post »
Post contributed by Ernst Hauber and Lida Fanara, Institute of Planetary Research, German Aerospace Center (DLR), Berlin, Germany.
Mars is an active planet, and several processes are currently shaping its surface. Among those, gravity-driven mass wasting produces surface changes that can be quantified in image data acquired before and after discrete events. As such changes are typically small in their spatial dimensions, the prime dataset to recognize them are pairs of HiRISE images (High Resolution Imaging Science Experiment; McEwen et al., 2007), with scales of ~25-50 cm/px. The manual identification of surface changes in these huge images (a single HiRISE image can have a size of several Gigabytes) is challenging, however, and requires significant efforts. In order to circumvent this massive demand on human resources while yet taking advantage of all images, automated methods need to be developed. Here we show an example of such methods which was applied to ice block falls at a steep cliff in Mars` north polar region.
Image 1: Block fall at a scarp on the north polar region of Mars near 83.796°N and 236.088°E. (a) «before» image acquired at 06 May 2014 (HiRISE image ESP_036453_2640). (b) «after» image acquired at 25 December 2019, showing a cluster of blocks that was displaced from the scarp to the right (east) (ESP_062866_2640). North is up in both images, scale bar is 100 m.
Post Contributed by Dr. Erin Leonard, Postdoctoral Fellow at the Jet Propulsion Laboratory, California Institute of Technology
Jupiter’s icy moon Europa has a geologically young surface (60-100 million years old), as evidenced by the sparsity of large impact craters. Studying the surface features on Europa allows insight into how resurfacing may have given it a youthful appearance. The majority of Europa’s young surface is made up of Ridged Plains terrain. This terrain has not been extensively studied before because it appears as a smooth and relatively bland in the global-scale images. However, in the few high-resolution images returned by the Galileo mission in the early 2000s, the Ridged Plains are revealed to consist of numerous ridges and troughs that have a range of morphologies—from crisscrossing each other in various directions to orderly sets of parallel structures (Image 1). But how did these ridges and troughs form?
Image 1: A variety of examples of ridged plains on Europa. Note the linear to curvilinear systematic ridge traces in all examples: (A) observation E4ESDRKMAT02 at 26 m/pixel, (B) observation 19ESRHADAM01 at 66 m/pixel, (C) observation 12ESWEDGE_02 at 29 m/pixel, and (D) observation 12ESMOTTLE02 at 16 m/pixel.
Post by Mallory Kinczyk, PhD candidate, Department of Marine, Earth, and Atmospheric Sciences, North Carolina State University
The formation of impact craters may be the most ubiquitous exogenic surface process in the Solar System. These craters take on many shapes and sizes and can hint at underlying rock types, tell us about the nature of the impactor, and can shed light on the body’s geological history. Even on bodies without atmospheres, erosive forces are at play, changing the crater shape through time via processes such as seismic shaking and disruption from debris thrown outward by subsequent, nearby impacts. Because Mercury is the only terrestrial planet without an atmosphere, it maintains a unique snapshot of the inner Solar System’s impactor population (Image 1) and, in turn, can shed light onto Earth’s own geological history.
Image 1: View of Mercury from the MESSENGER spacecraft, which orbited Mercury between 2011 and 2015 (Image PIA17280). A variety of impact crater sizes and shapes are evident from very fresh craters to subdued to almost completely obliterated crater forms. Bach crater (arrow) hosts a well-defined central peak ring, but its subdued form indicates that it has been disrupted by subsequent craters. Image Credit: NASA/Johns Hopkins University Applied Physics Laboratory/Carnegie Institution of Washington.
Geomorphologists working with Mars share a frustration of not being able to visit their objects of investigation. To counter this, a commonly used approach is to look for environments on Earth that resemble those studied on Mars. This approach, called Earth-analogue studies, helps to guide our line of reasoning in deciphering formation mechanisms of specific martian landforms of interest. Mars, being the most earth-like planet in the solar system hosts numerous landscapes and landforms that in plan-view show remarkable similarities to known features on Earth. Especially striking examples are martian glacial flow-like features and gullies to that resemble terrestrial glaciers and fluvially-incised ravines, respectively. As a consequence their Earth counterparts have been studied with great intensity for the last couple of decades. Although correspondences in form may guide our way of thinking of plausible formative processes by reference to Earth, the approach is not without pitfalls. For example, experimental studies in Mars climate chambers have shown that fluvially triggered slope processes may be of a completely different nature under Mars’ atmospheric conditions of low pressure combined with low temperatures, but the resulting landform looks about the same. This is a problem of equifinality (i.e. convergence of form), something that also terrestrial scientists encounter but which is a major challenge in planetary geomorphology (e.g. Hauber et al. 2011; Zimbelman 2001). One way to try to minimize equifinality is by taking whole landform assemblages into account where different types of landforms may have some genetic linkages.
Post by Dr. A Lucas, CNRS Research Scientist, Université de Paris, Institut de physique du globe de Paris, CNRS, F-75005, Paris, France
Titan, Saturn’s largest satellite, has proved to be a world that is both strange and yet so familiar to us. Mountains, lakes, drainage systems and dune fields (Images 1-2) cover its surface. Methane on Titan occupies a similar position to water on Earth. It participates in climatic cycles. Moreover, its photodissociation in the upper atmosphere is responsible for the soot rains that fall on the surface of this icy world. The fate of these grains composed of organic materials is just as essential. Indeed, winds sometimes mobilize them. Over long, very long periods of time, this granular transport is responsible for the formation of vast dune fields located at the equator. But after 13 years of exploration by the Cassini probe, these dunes have not revealed all their secrets. In particular, their morphodynamics are widely debated. Are these bedforms remains from an old time, are they still active today? What is their growth rate? And what is their resulting sediment flux?
Image 1: Despeckled T8 swath SAR image over the Belet sand dunes located at the Equator of Titan. The dark longitudinal features are the micro-wave absorbent dunes composed of sand made of organics molecules. The bright areas are rough topographic reliefs revealing the icy bedrock beneath the organic sediment cover. Glints (bright spots) are detectable of the crest of some dunes due to specular reflection on their avalanche side.
Post by Dr. Norbert Schorghofer. Senior Scientist, Planetary Science Institute
Recurring slope lineae (RSL) are dark narrow streaks on Mars that have puzzled scientists since their discovery in 2011. Image 1 shows a 3-dimensional perspective of a landscape with some of these flow features. RSL form and grow annually and mostly in the warm season, so the mechanism by which they form and grow is tantalizing. To what extent are RSL related to temperature or water? In rugged terrain there are stark temperature contrasts between pole‐facing and equator‐facing slopes that infrared cameras on Mars-orbiting spacecraft cannot spatially resolve. New modeling capabilities make it possible to overcome this limitation and provide surface temperatures at high spatial resolution.
Image 1: An image combining orbital imagery with 3-D modeling shows RSL on a slope inside Newton Crater. Image credit: NASA/JPL-Caltech/University of Arizona.
Post by Dr. O. Ruesch, European Space and Technology Center, European Space Agency (ESA), the Netherlands.
A fragmented block is referred to a cluster of fragments formed by the disruption of a parent block. The identification of such features on planetary surfaces is possible due to the minor spatial dispersion of the fragments away from the parent block. This morphology is to be distinguished from clusters of fragments formed by mass wasting like rockfall or disintegration during block rolldown. Observations of fragmented blocks have been reported on almost every rocky planetary body where images captured by orbital and surface craft resolved features in sufficient spatial detail. Despite the fact that disrupted blocks can reveal important clues on the formation process of soil (regolith) on planetary surfaces, they have started to receive attention only in recent years.
On the airless surface of the Moon, fragmented blocks display a wide range of morphologies (Images 1 and 2). In general, the configurations of the fragments can be described by a continuum from highly catastrophic to sub-catastrophic. Image 1 shows an example of a catastrophic fragmentation where the number and size of the fragments indicate that the parent block was much larger than the largest fragment. The radial pattern formed by small fragments and brighter areas is diagnostic of disruption by a meteoroid impact.
Image 1. Example of a block fragmented catastrophically near crater Copernicus on the Moon, where the largest fragmented in considerable smaller than the original parent block. LROC/NAC image M127063668LE. http://bit.ly/2mAl0CB
Post by Dr. Richard Soare, Geography Department, Dawson College, Montreal, QC H3Z 1A4, Canada.
On Earth, hydrostatic or closed-system pingos (CSPs) are perennial ice-cored mounds formed by the freeze-thaw cycling of water when or as thermokarst lakes lose their water by drainage or evaporation. The mounds vary in shape from circular or sub-circular to elongate, are sub-kilometer in their long axes and may reach decameters in height. Some mounds show summit depressions, radial fractures and tiers. As the CSPs degrade, small-scale debris flows or slumps may occur; end-stage degradation often is marked by debris-laden ramparts elevated symmetrically or asymmetrically above the surrounding terrain. The presence of closed-system pingos on Mars might thus inform us about the planets’ past and/or present hydrologic conditions (Image 1).
Image 1: Candidate closed-system pingos (CSPs) at the mid latitudes of Utopia Planitia, Mars (HiRISE image ESP 027650_2275). Note the location of the potentially ice-cored mounds in polygonised thermokarst-like depressions and the interesting mound morphologies. On Earth, CSPs form almost uniquely in the midst of drained or water-depleted thermokarst lakes (alases) that are nested in ice-rich terrain. On Mars, we propose that the depressions form by the sublimation-driven loss of near-surface ice; this would also be the principal driver of ice-core depletion and/or collapse. Mounds 1-4 exhibit circular shapes and summit depressions that are commonplace amongst CSPs in regions like the Tuktoyaktuk Coastlands of northern Canada. The crescentic shape of Mound 5 could be a marker of mound collapse and the deflation/sublimation of a near-surface ice core, the geological buttress of pingo topography on Earth.
Post by Dr. Stéphane Le Mouélic, Laboratoire de Planétologie et Géodynamique, CNRS UMR6112- University of Nantes, Nantes, France.
Titan is one of the most fascinating bodies of our Solar System. Bigger than Mercury, this satellite of Saturn is veiled by a thick atmosphere of nitrogen containing a few percent of methane. Aerosols formed in the atmosphere by a complex chemistry triggered by the solar UV irradiation produce a global haze totally masking the surface to the naked eye. During 13 years, from July 2004 to September 2017, the Cassini spacecraft orbited Saturn. It took advantage of gravity assist maneuvers to perform 127 equatorial and polar flybys of Titan. Data from the Visual and Infrared Mapping Spectrometer (VIMS) onboard Cassini revealed the distribution of the main compositional units of the surface of Titan (Image 1). The inset in Image 1 shows the 84 km-diameter Selk crater, one of the primary targets chosen for the next New Frontier “Dragonfly” mission, a mobile rotorcraft-lander planned to be launched in 2026.
Image 1: False color composite of Titan with the red controlled by the 1.59/1.27 µm, green by the 2.03/1.27 µm and blue by the 1.27/1.08 µm band ratios. The equatorial dune fields appear in a consistent brown color. Selk crater is shown in the inset. Credits NASA/JPL/Univ. Arizona/CNRS/LPG.
Post by Valentin Bickel, PhD student, Department of Earth Sciences, ETH Zurich, CH & Department Planets and Comets, Max Planck Institute for Solar System Research, GER.
One of the most intriguing objects on the surface of the Moon are the “rolling stones”, also known as lunar rockfalls or rolling boulders (Image 1). These boulders are abundant all over the Moon and have sizes that range from a couple of meters to several 10s of meters. Lunar boulders are believed to be displaced by moonquakes or impacts and can carve tracks with lengths that range from a couple of meters to several kilometers (Image 1; Xiao et al., 2013; Kumar et al., 2016). Besides their value for geomorphological analyses, these boulder tracks provide insights into the mechanical behavior and the trafficability of the lunar “soil”, the regolith (Bickel et al., 2019).
Image 1: A number of large and small boulders with tracks at the bottom of a lunar slope. The analysis of tracks provides insights about the mechanical properties of the regolith and is performed using high-resolution satellite imagery, taken by NASA’s Lunar Reconnaissance Orbiter Narrow Angle Camera (NAC). Detail of NAC Image M113934119LC.
Post by Dr. F. Salese, Marie Curie Postdoctoral Fellow, Faculty of Geoscience, Utrecht University.
Groundwater had a greater role in shaping the Martian surface and may have sheltered primitive life forms as the planet started drying up. Observations in the northern hemisphere show evidence of a planet‐wide groundwater system. The elevations of these water‐related morphologies in all studied basins lie within the same narrow range of depths below Mars datum (Image 1) and notably coincide with the elevation of some ocean shorelines proposed by previous authors. Most previous studies on Mars relevant groundwater have proposed models, but few have looked at the geological evidence of groundwater upwelling in deep closed basins in the northern hemisphere equatorial region. Geological evidence of groundwater upwelling in these deep basins is a key point that will help to validate present-day models and to better constraint them in the future.
Image 1: Morphologies inside several basins. a) Crater #15 shows the presence at the same time of delta, sapping valleys, debris and hummocky terrain. The basin floor is flat. b) Crater #12 shows stepped delta, terraces, shorelines and flow structures at about the same topographic elevations. c) Sapping valley and related stepped delta in crater #18. d) Sapping valley and related stepped delta along with fan and exhumed channels in crater #12. e) Crater #16 shows well-preserved outcrops of debris flow. f) Sapping valley with related delta at -4100m inside crater #22.
Post by Giulia Magnarini, PhD candidate, Department of Earth Sciences, University College London, UK.
The availability of high resolution imagery of the surface of Mars from NASA’s Mars Reconnaissance Orbiter CTX and HiRISE cameras (NASA PDS) allow us to reconstruct fantastic 3D views of the martian topography using stereophotogrammetry technique. Digital terrain models (DTMs) are obtained using the difference in two images of the same target taken from different angles. In the process, orthoimages are generated and draped over the DTM. CTX stereo-derived DTMs have 20 m/px resolution; HiRISE stereo-derived DTMs have 1-2 m/px resolution. This technique is applied to the study of martian long runout landslides and it represents a powerful tool, as the 3D reconstruction allows detailed observations and morphometric analysis of these landforms and their morphological features (Images 1-3).
Image 1: Long runout landslide in Ganges Chasma, Valles Marineris, Mars. CTX stereo-derived DTM at the Mineral and Planetary Sciences division of the Natural History Museum in London. Vertical exaggeration 2x. Image pair: P20_008681_1722_XN_07S044W and P20_009037_1718_XN_08S044W.
Post by Joshua E. Hedgepeth, PhD student, Centre for Planetary Science and Exploration, University of Western Ontario, Canada.
Discovered by Huygens in 1656 (Huygens, 1656), the surface of Titan was obscured by its atmosphere for centuries (Campbell, 2003; Smith et al., 1996; Muhleman et al., 1995, 1990; Kuiper, 1944). In 2004, we finally obtained high resolution images of the surface with Cassini RADAR. Cassini was equipped with a Ku-band (2.17 cm λ) radar instrument with 5 beams for collecting data (Elachi et al., 2004; Stofan et al., 2012). The long wavelength band was able to penetrate Titan’s thick haze to perform radiometry, scatterometry, altimetry and synthetic aperture radar (SAR) imaging of the surface. The SAR mode captured the highest resolution images of the surface of Titan, as high as 175 meters per pixel (Elachi et al., 2004; Lopes et al., 2010). While this may not be as high resolution as the images we have of other moons, it is high enough to obtain some spectacular images. In this blog we have a detailed look at the craters on the surface of Titan.
Image 1: Selk crater (D=84 km) located at 199.1, 6.9 latitude and longitude. The dark crater floor is representative of the smooth material, and around it is the bright rough ejecta material.
Planetary Geomorphology Image of the Month 