# The role of liquid water in recent surface processes on Mars

Susan J. Conway, David Stillman

► To cite this version:

Susan J. Conway, David Stillman. The role of liquid water in recent surface processes on Mars. Geological Enigmas, Elsevier, pp.207-261, 2021, 10.1016/B978-0-12-820245-6.00009-4. hal-03595959

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The role of liquid water in recent surface processes on Mars

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3 Susan J. Conway¹*

4 David E. Stillman²

5 1. CNRS UMR 6112 Laboratoire de Planétologie et Géodynamique, Université de Nantes, France.

6 2. Dept. of Space Studies, Southwest Research Institute, 1050 Walnut St. #300, Boulder, CO 80302, USA
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Abstract

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We provide an up-to-date review on the different landscapes and landforms that have been attributed to the action of liquid water in the Amazonian epoch on Mars and the current state-of-the-art regarding their interpretation. This chapter accompanies Chapter DUNDAS (pages xxx-xxx) where the counterarguments are presented. The Amazonian epoch is thought to be dominated by hyper-arid climate conditions hostile to surface liquid water, and our review reveals that this steady-state is likely to be punctuated by episodic appearances of liquid water at the surface. The proposed sources of liquid water are varied: groundwater, thawing of surface-ice, deliquescence, ground-ice or glaciers, with triggers as variable as microclimates, climate-shifts, geothermal anomalies and impact cratering.

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Our review covers recently active surface processes in the form of Slope Streaks, dark dune flows and Recurring Slope Lineae, all hypothesised at one point or another to be seeps of liquid water. We then cover landscapes and landforms which are proposed to be a result freeze-thaw cycles in the recent past, including gullies, lobate forms on hillslopes, pingo-like mounds, low-centred polygons, patterned ground and ice-loss landscapes. We present evidence that liquid water has been produced at the base of glacial landforms during the Amazonian, resulting in sinuous ridges (eskers) and enhanced crater-wall erosion. Liquid water brought up-from depth is thought to have produced mud-volcanos and other features related to sedimentary volcanism in the Amazonian and small fluvial channels amongst other features are thought to be related to melting induced by periodic impact events.

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We end by summarising the importance of the search evidence of liquid water on Mars and by proposing solutions to the current impasses where progress is hindered due to limitations in data or our understanding.
## 1 Present-day and recent surface conditions on Mars

Of the planets of the Solar System Mars and Earth are the only two inner planets whose surfaces are thought to have hosted abundant liquid water early in their history (4.5-3.5 Ga). Liquid water is key for supporting life and seeking it out is a key tenet of astrobiology (Abrevaya et al. 2016). Hence, when searching for signs of extra-terrestrial life, space agencies have made Mars one of their principal targets. Mars is a rocky planet located outside Earth's orbit at a solar distance between 1.38 AU and 1.67 AU and has 1/3 the radius of the Earth (the mean martian radius is 3389.5 km). It orbits the sun in 687 earth days (~2 Earth years) and has a very similar day-length to the Earth (24hrs 37min). Its present-day axial tilt is close to the Earth's (25.19° compared to 23.44° on the Earth), giving rise to similar temporal and spatial patterns of insolation over a martian year. The eccentricity of the martian orbit is greater than the Earth (0.0935 compared to 0.0167), and the precession of its aphelion/perihelion through the seasons leads to greater seasonal amplitudes and asymmetries than experienced on Earth. For example, with the current position of perihelion, the southern winters are colder and shorter, compared to the southern winters which are warmer and longer. Mars lost its internally generated magnetic field approximately 3.6 Ga (Lillis et al. 2008; Milbury et al. 2012) that probably led to the loss of the majority of its atmosphere (Jakosky & Phillips 2001). Surface geomorphology and sedimentary composition both point to declining activity of liquid water on Mars' surface over the last 3.5 Ga, which is strongly linked to the loss of Mars' atmosphere (Jakosky & Phillips 2001). Mars' present-day atmospheric pressure ranges from 5-12 mbar (e.g., Hess et al. 1980; Ordonez-Etxeberria et al. 2019) and its main constituent gas is CO₂.

The low atmospheric pressure combined with the greater solar distance, mean that surface temperatures on Mars are on average much lower than on Earth (global mean is ~210K or -60°C) and the diurnal variation of temperature much greater. The lowest surface temperature is buffered around the condensation temperature of CO₂ at around 150K, while peak summer temperatures on and directly above favourable surface materials and orientations can exceed 300K (e.g., Stillman et al. 2014). The pressure and temperature conditions on present-day Mars mean that water should only be stable in solid or gaseous state. Atmospheric humidity on Mars is measured in precipitable microns, i.e. if the whole atmospheric column was condensed, only microns of water would result. Humidity typically ranges up to hundreds of precipitable microns (Smith 2002), which is enough to generate, thin, yet visible surface frosts (Svitek & Murray 1990). These general ranges of atmospheric pressure, humidity and temperature are below the triple point for liquid water, making this phase unlikely to occur in bulk.

However, there is theoretical and indirect evidence to support the possibility of thin films of water or brines developing within the martian soil at the present-day (e.g., Boxe et al. 2012; Jouglet et al. 2007; Milliken et al. 2007; Kereszturi & Rivera-Valentin 2012; Möhlmann 2008) or in the recent past (Sizemore et al. 2015). Pores within a regolith where water ice is also present can form microenvironments where the sublimation of the water ice can lead to increased humidity conditions (beyond those of the atmosphere). Under such conditions even at temperatures below the frost point transient liquid water films can form at the regolith or ice grain contacts and surfaces. Since data sent by the Phoenix mission (Hecht et al. 2009) perchlorate salts are now thought to be widespread on Mars (Navarro-González et al. 2010; Glavin et al. 2013a) and depress the freezing point and evaporation rate of water on Mars favouring the formation of thin water films. Atmospheric humidity can increase to levels and beyond near the surface within the diurnal cycle (e.g., Maltagliati et al. 2011) and where the soil is in diffusive exchange with the atmosphere this could promote the formation of thin films of brine/water by deliquescence (Pál & Kereszturi 2017; Cull et al. 2014). Additionally, even before a water saturated
atmosphere is reached, perchlorate salts can adsorb water from the atmosphere via deliquesce and can form brines, which then have remarkable metastable properties (e.g., Gough et al. 2011; Nuding et al. 2015; Toner et al. 2014b; Primm 2018; Primm et al. 2019; Primm et al. 2018)(e.g., Gough et al., 2011; Nudging et al., 2014; Toner et al., 2014; Primm et al., 2018; Primm and Stillman, 2019). Whether or not such films constitute habitable environments has not reached consensus, but their water activity tends to be too low to support life (Martín-Torres et al. 2015; Jones & Lineweaver 2012).

Some researchers have theorised that bulk transient liquid water could exist in the recent martian climate, either brought-up from underground aquifers (e.g., Gaidos 2001; Goldspiel & Squyres 2011; Malin & Edgett 2000; Stillman et al. 2016) or melted from ice under extremely rare favourable combinations of temperature, and humidity conditions (e.g., Chevrier & Rivera-Valentin 2012; Grimm et al. 2014; Hecht 2002; Kossacki & Markiewicz 2010). Melting from surficial ice is particularly difficult as there is a kinetic barrier to melting even if conditions are conducive: because latent heat is consumed by sublimation at the surface of the ice which prevents the temperature from rising sufficiently to allow melting to occur (Mellon & Phillips 2001).

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89 Many researchers have hypothesised that Mars' climate was substantially different in the recent past
90 and may have been more conducive than at the present-day to the emergence and stability of liquid
91 water (e.g., Costard et al. 2002; Richardson & Mischna 2005). This is because Mars' orbital parameters
92 (Fig. 1) are much more variable than those of Earth and we already know that Earth's climate undergoes
93 glacial-interglacial cycles as a result of periodic changes in orbital obliquity, eccentricity and the solar
94 longitude Ls of perihelion – so called Milankovitch cycles – which change the amount and distribution of
95 incoming solar radiation. Hence, the larger variations experienced by Mars should have more impact on
96 the climate system than on Earth. Having a lower mass (i.e. moment of inertia) than the Earth and
97 without a large Moon to stabilise the planet, Mars' axial tilt (obliquity) has varied between 15° and 40° in
98 the last 5 Ma (compared to ±1.3° around the present value for Earth) with a period of 120 ka. Mars'
99 eccentricity varies between 0 and 0.12 (compared to 0 to 0.068 for Earth) over 95 ka and Ls of perihelion
100 takes ~51 ka years to return to its starting season (Fig. 1) (Ward 1992; Ward 1979; Ward 1974). Global
101 Climate Models (GCMs) have been used to try and predict the influence that these orbital changes
102 should have on Mars' climate. Despite these large variations in incoming radiation, the low atmospheric
103 pressure on Mars means that reaching the triple point of water is still challenging (Richardson & Mischna
104 2005). Using current values for atmospheric pressure with obliquity at 35° Richardson and Mischna
105 (2005) predict at most 100 days per year of transient liquid water (i.e. the time above the triple point)
106 located at the base of Hellas Basin. It is hypothesised that at very high obliquities (>35°) sufficiently
107 intense incoming solar radiation, at certain latitudes and times of year, means that ground ice could melt
108 if it is protected from the low atmospheric pressure by an overburden slowing heat-loss by sublimation
109 and impeding gas diffusion (e.g., Williams et al. 2009). At such obliquities, the two water ice polar caps
110 are sublimated away due to the increased insolation at the poles. This significantly increases the
111 atmospheric water vapour concentration (Mellon & Jakosky 1995) and leads to a massive redistribution
112 of water ice by cold trapping in the mid latitude region that is the coldest. It is also the region that
113 experiences the most intense increases in solar radiation, so it is at the mid-latitudes where the most
114 intense surface-atmosphere exchange of water is thought to occur. Conversely, at obliquities <17°, the
115 atmospheric water vapour concentration decreases by more than an order of magnitude from the
116 current value (Mellon & Jakosky 1995) and prevents migration of any atmospheric water vapour. The
117 recent discovery of large volumes of CO2 ice sequestered in the south polar cap (Bierson et al. 2016;
118 Manning et al. 2019; Phillips et al. 2011), opens up the possibility for a slightly higher atmospheric
119 pressure (Buhler et al. 2019; Manning et al. 2006) than assumed by previous modelling studies, which
120 would allow liquid water to be stable for a few more degrees above freezing before it boils.
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121 These conditions – that of a hyper-arid cold desert – have thought to have reigned on Mars during the
122 whole Amazonian period (since ~3 Ga) even if uncertainties in the orbital solutions do not allow us to
123 predict the orbital parameters further than several tens of millions of years in a reliable way.

124 2 Present-day reservoirs of water on Mars

125 Even if surface conditions allow for the presence of liquid water, there must be a source of water (ice or
126 vapour) for liquid water to be produced. Precipitation in the form of rain is generally accepted to be
127 highly unlikely during the Amazonian (since ~3 Ga). There are three potential reservoirs for water:
seasonally deposited ices/frosts from vapour in the atmosphere, surface ice reservoirs, or underground aquifers. Here, surface ice reservoirs are defined as those ices which can be reasonably considered to have exchanged with the atmosphere under Amazonian climate conditions. We discuss each of these sources in turn below and their likelihood for generating liquid water.

## 2.1 Seasonal ices and frosts

The most visible manifestation of Mars' seasonal cycle is the waxing and waning of the seasonal polar caps, which were noted even in early telescopic observations (Herschel 1784). The seasonal polar caps comprise atmospherically deposited seasonal ices that brighten the surface. The majority of the ice is comprised of CO₂ and as a result the deposition of the seasonal ices causes a substantial fluctuation in the atmospheric pressure (Haberle et al. 2017). A minor component of water ice is included in these ices (Appéré et al. 2011; Brown et al. 2012; Brown et al. 2010). The ices are thought to originate as directly condensed frosts or fall as snow (e.g., Giuranna et al. 2008; Hayne et al. 2014), with direct observation being hindered by the polar hood (low lighting combined with a hazy atmosphere). The continuous seasonal polar caps, comprised of CO₂ and a minor component of H₂O, extend to 50-55° in both hemispheres (e.g., Piqueux et al. 2015) and are thought to reach a thickness of metres near the poles (e.g., Karatekin et al. 2006; Smith et al. 2001). Discontinuous deposits of ices are found down to latitudes of 30° in both hemispheres (Dundas et al. 2019b; Vincendon et al. 2010a; Vincendon et al. 2010b), where water ice is found to lower latitudes than CO₂ ice. Discontinuous CO₂ extends to 35°S on steep pole-facing slopes in the southern hemisphere (Vincendon et al. 2010a), but its extent has not yet been measured in the northern hemisphere. Sublimation of the seasonal ices drives peaks in atmospheric water vapour to around tens of precipitable microns (Pankine et al. 2010; Pankine & Tamp Parisi 2019). Water ice frost forms an annulus around the retreating seasonal ices and as it sublimates is cold-trapped onto the retreating CO₂ ice before sublimating again later. Exchange of water vapour between hemispheres is strongest during the solstices (Haberle et al. 2017).

In general, water frosts and ices deposited seasonally are deemed to be unlikely sources of liquid water at the present-day and in the recent past, due to the very small quantities and the kinetic barrier to melting discussed in Section 1.

## 2.2 Surface Ice reservoirs

Surface ice reservoirs are our most accessible evidence that Mars has abundant water, and these can be split into:

158 - The polar caps (~$10^6$ km³ each) (Plaut et al. 2007; Selvans et al. 2010)

159 - Mid-latitude glaciers (~$10^5$ km³) (Levy et al. 2014)

160 - Ground ice (~$10^4$ km³) (Conway & Balme 2014; Kreslavsky & Head 2002; Mustard et al. 2001)

Radar data have confirmed that the polar caps are almost exclusively water ice (Plaut et al. 2007; Selvans et al. 2010). The polar caps are the most voluminous reservoir of water on Mars and form upstanding broadly domal landforms with up to three kilometres of relief (Plaut et al. 2007; Selvans et al. 2010). The internal structure of the two caps is broadly similar, comprising layers of ice with variable dust-contents, visible in surface exposure at internal scarps and at depth via radar sounding (e.g., Becerra et al. 2017; Blasius et al. 1982; Fishbaugh & Hvidberg 2006; Plaut et al. 2007; Selvans et al. 2010) and thus are generally accepted to form via atmospheric deposition of water ice. Both caps are dissected by spiral troughs which are thought to form during cap accumulation through a balance between sublimation and
169 accumulation (Bramson et al. 2019; Smith et al. 2013; Smith & Holt 2010) and they are more marked in
170 the northern cap than in the southern cap. However, in the detail the two caps are quite different, likely
171 as a result of the huge difference in altitude between the two hemispheres (base at -5000 m in the north
172 and +1000 m in the south), which has a profound effect on atmospheric circulation (e.g., Haberle et al.
173 2017). The northern polar cap has a younger surface age than the southern cap (a few thousand
174 compared to ~10 Ma; Herkenhoff & Plaut 2000; Landis et al. 2016), and its bulk is estimated to have
175 formed in the last few millions of years (Jakosky et al. 1995; Laskar et al. 2002; Levrard et al. 2007) even
176 though its surface may be being continuously renewed (Herkenhoff & Plaut 2000; Landis et al. 2016). It is
177 thought to comprise a lower unit of ice-rich sands and an upper unit of polar layered deposits. It is
178 broadly centred on the pole and is surrounded by outliers which formed separately (Bapst et al. 2018;
179 Brothers & Holt 2016; Brown et al. 2008; Conway et al. 2012). Hovius et al. (2008) suggested that some
180 re-entrants into the cap represented channels carved by water released by volcanic heating, but later
181 work showed these features were more consistent with wind erosion (Brothers et al. 2013). In contrast,
182 the southern polar cap is located off-centre with respect to the southern pole and only a very small
183 proportion of the cap has exposed water ice at the surface in summer (Hansen et al. 2005), compared to
184 the complete exposure expressed by the northern cap.
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Gully Orientation

Equatorward

E-W

Poleward

No preference

Water-equivalent
hydrogen in an lower
layer (%)

< 5

5 - 10

10 - 15

15 - 50

>50

Small-scale lobes

Elevation (m)

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Viscous Flow
Features
Figure 2: Global distribution of relevant landforms in the quest for evidence of liquid water on Mars. (a) Mars Odyssey Thermal Emission Spectrometer global albedo map overlain by the global distribution of Recurring Slope Lineae (RSL), where the legend to the right shows how many RSL attributes are displayed at each site, a confirmed RSL site being in red having evidence of incremental lengthening, recurrence and fading. (b) Mars Orbiter Laser Altimeter (MOLA) derived shaded relief map overlain with the locations of gullies as published by Harrison et al. (2015). (c) Semi-transparent MOLA-derived shaded relief map with water equivalent hydrogen data from Pathare et al. (2018) overlain with locations of small-scale lobes from Johnsson et al. (2018). (d) Semi-transparent MOLA-derived shaded relief map with colour-keyed MOLA-elevation overlain by the locations of glacial forms, or VFF as mapped by Levy et al. (2014).

The bright part of the cap is instead CO₂ ice. The deep structure of the polar cap has been harder to decode than for the northern cap because there are significantly fewer internal radar reflections. The cap is believed to have a series of packets of polar layered deposits with significant zones where CO₂ is thought to be sequestered (Bierson et al. 2016; Manning et al. 2019; Phillips et al. 2011). Much of the cap is covered by dust and hence has a lower albedo than exposed ice (Herkenhoff 2001). The surface of the exposed cap is covered by a metres-thick layer of CO₂ ice which is currently retreating at metres per year (Buhler et al. 2017; Byrne 2003; Thomas et al. 2009). There are numerous outliers of the southern cap and these might represent a more extensive cap in the past (Sori et al. 2019; Westbrook 2009). Ancient units that surround the south polar cap (e.g., Fastook et al. 2012; Scott & Tanaka 1987; Head & Pratt 2001) are thought to represent water production at the base of the cap, which will be discussed in more detail in Section 3.5 in the context of a recent report of sub-glacial lakes found via radar sounding under the southern polar cap (Orosei et al. 2018; Lauro et al. 2020).

The mid-latitude glaciers are common and widespread across areas with topographic relief in the mid-latitudes (30-60°; Fig. 2d) and have been classified by morphology and context into several groups including: viscous flow features (VFF), Concentric Crater Fill (CCF), Lobate Debris Aprons (LDA), Lineated Valley Fill (LVF) and Glacier-Like Forms (GLF). The glacial origin of all these features has now been generally accepted with the most convincing evidence coming from orbital sounding radar (Holt et al. 2008; Petersen et al. 2018; Plaut et al. 2009) and further conformed by their overall shape/form and surface textures (e.g., Baker & Carter 2019a; Baker & Carter 2019b; Head et al. 2010; Karlsson et al. 2015; Levy et al. 2009c; Mangold 2003; Parsons et al. 2011). These glaciers are thought to comprise almost pure ice with a metres- to decametres-thick debris cover derived from the headwall and englacial materials. The present-day systems are believed to be out of equilibrium with the current climate and show extensive evidence for retreat and downwasting (Brough et al. 2016). Superposition relationships (Hepburn et al. 2019) and estimations of age derived from crater-size frequency distributions on their surfaces suggest that they formed during several "glacial maxima" over the last tens to hundreds of millions of years (Baker et al. 2010; Berman et al. 2015; Hartmann & Werner 2010; Morgan et al. 2009). Each individual feature is likely to represent a series of ice-accumulation events (Fastook & Head 2014; Weitz et al. 2018). As a general rule, martian glaciers are thought to be cold-based, i.e. no melting occurs near the base as is common in terrestrial glaciers, and they move via viscous deformation of the ice, rather than sliding over their bed (e.g., Parsons et al. 2011), which is consistent with their form. Possible exceptions to this general rule will be discussed in Section 3.5.

Ground ice is extensive on Mars and is spatially associated with surface deposits that tend to drape the topography, variously termed "mantling units", "mantle", latitude dependant mantle", "pasted on terrain", "dust-ice mantle" amongst other terms. The main line of evidence for ground ice comes from
the neutron spectrometer on Mars Odyssey (Boynton et al. 2002; Feldman et al. 2011; Feldman et al. 2008; Feldman et al. 2007; Feldman et al. 2004; Maurice et al. 2011; Pathare et al. 2018; Wilson et al. 2018), where the content of ice in the top metre of the regolith is estimated to be 50-80% at latitudes >50°N/S (Fig. 2c). These predictions were shown to be locally valid by the Phoenix lander, which found ice at depths of a few centimetres both in the form of pore-filling ice and excess ice (Mellon et al. 2009).
Surface temperature data from Mars Climate Sounder and Thermal Emission Imaging System indicate that ice occurs within centimetres to tens of centimetres of the surface at latitudes down to 35°N/45°S (Piqueux et al. 2019). In addition, high ice-contents from the neutron data correlation spatially with the observations of polygonally patterned ground (e.g., Levy et al. 2010b; Mangold 2005), which is generally accepted to be a result of thermal contraction of ice-cemented ground (Mellon 1997). Phoenix landed amongst such polygons and excavated and detected ground ice with a mean depth of 4.6 cm (Mellon et al. 2009). Additionally, Phoenix also detected segregated ice above the ground ice table. New impacts that reveal ice in their ejecta generally confirm this distribution (Byrne et al. 2009; Dundas et al. 2020; Dundas et al. 2014), and even extend to lower latitudes (~40°) where the ice is predicted to be deeper than can be detected by the neutron spectrometer. At these latitudes, the surface mantling units become dissected, exposing layering and exhibiting pitting (Milliken et al. 2003; Schon et al. 2009b). These discontinuous mantles are found preferentially on pole-facing slopes at lower latitudes (e.g., Christensen 2003). This latitude marks the boundary where the topography at scales of hundreds of metres becomes measurably smoother where these mantles are present (Kreslavsky & Head 2002; Kreslavsky & Head 2000).

Ground ice is thought to be responsive to climate change and in diffusive equilibrium with the current climate (Aharonson & Schorghofer 2006; Mellon & Jakosky 1993; Schorghofer & Aharonson 2005). The relationship between the near-surface ground ice detected by the neutron and thermal orbital data and the surface mantling units (with their ice content) is ambiguous. In some cases such surface mantling units can be locally decametres thick and have evidence for massive ice, from geomorphic and radar evidence (Bramson et al. 2015; Conway & Balme 2014; Dundas et al. 2018; Stuurman et al. 2016). Such enrichments cannot come about from atmospheric diffusion into the regolith because of self-blocking, even if this can be bypassed by contraction cracking (Fisher 2005) only tens of centimetres depths can be reached (see Section 3.4.5). The alternative is that these mantles represent units laid down under past high-obliquity climate conditions when snow/frost could be directly accumulated at the surface (Head et al. 2003; Madeleine et al. 2014; Madeleine et al. 2009) and protected from later sublimation by the accumulation of a sublimation lag of atmospherically derived dust centimetres in thickness. However, a universal composition of nearly pure ice for these units is not supported by the geomorphic evidence, which rather suggests development of the ground ice within a regolith matrix, i.e. excess ice (e.g., Levy et al. 2010b; Pathare et al. 2018; Soare et al. 2015). We discuss the development of excess ice further in the context of ice-loss landforms in Section 3.4.5.

Ground ice is generally considered to be an improbable, but still viable source for liquid water in the present and recent-past, as the soil above protects it from sublimation, removing some of the constraint of the kinetic barrier to melting.

## 2.3 Subsurface water reservoirs

Numerous authors have hypothesised the presence of aquifers on Mars (e.g., Abotalib & Heggy 2019; De Toffoli et al. 2019; Gaidos 2001; Goldspiel & Squyres 2011; Salese et al. 2019), initially as a way to explain the fate of the large quantity of liquid water that is thought to have existed on Mars early in its history,
272 as evidenced by the existence of the extensive valley networks (e.g., Craddock & Howard 2002; Hynek & Phillips 2003; Kereszturi & Petrik 2020; Luo et al. 2017; Penido et al. 2013; Williams & Phillips 2001; Hynek et al. 2010), the clay deposits found in Noachian terrains (e.g., Carter et al. 2015; Ehlmann et al. 2011b; Loizeau et al. 2007; Lowe et al. 2020; Lowe et al. 2020; Milliken & Bish 2010; Poulet et al. 2005), the outflow channels dated to the Hesperian or later (Andrews-Hanna & Phillips 2007; Komar 1979; Leask et al. 2007; Marra et al. 2014; Max & Clifford 2001; Montgomery & Gillespie 2005; Vijayan & Sinha 2017; Wilson 2004) and the hypothesised northern ocean (e.g., Baker et al. 1991; Carr & Head 2019; Clifford & Parker 2001; Palumbo & Head 2019). Mars is thought to possess a deep cryosphere (Carr 1996; Clifford et al. 2010; Grimm et al. 2017; Lasue et al. 2019), where this cryosphere is impinged on by the geothermal heat flux, liquid water is thought to exist. Tectonic and surface temperature fluctuations perturbing the pressure state of this aquifer are thought to be responsible for the hundred-kilometre scale outflow channels which occurred in the Hesperian (Baker 1979; Carr 1979; Coleman 2005; Harrison & Grimm 2008; Marra et al. 2014). Smaller, yet similar perturbations have been proposed to explain gullies and RSL (Abotalib & Heggy 2019; Gaidos 2001), bringing up water from kilometres in depth. Shallower surface aquifers, including perched aquifers, have also been proposed (Goldspiel & Squyres 2011; Marquez et al. 2005), but these are within reach of the orbital ground penetrating radars, such as the European Space Agency's MARSIS and NASA's ShaRAD, which have found no evidence for such bodies of water (Nunes et al. 2010). The non-detection of shallow aquifers by orbital radar is not definitive proof of their absence because the penetration of these radars into the martian crust is much lower than initially anticipated (Stillman & Grimm 2011) and they are limited to detecting interfaces that aquifers may in fact not present (Farrell et al. 2009; Nunes et al. 2010).

293 Circulation of groundwater in the past has been hypothesised from in situ observation of post-digenetic alteration products, such as veins and nodules (e.g., Chan et al. 2004; L'Haridon et al. 2018; Stack et al. 2014) found within the sedimentary rocks deposited in impact craters and evidence from meteorite studies (Bridges et al. 2001; Gillet et al. 2002; Schwenzer et al. 2016), evidence for sedimentary volcanism (e.g., Brož et al. 2019; Gallagher et al. 2018; Oehler & Allen 2010; Okubo 2016; Skinner & Mazzini 2009; Wheatley et al. 2019), and orbital detection of hydrothermal alteration products (e.g., Bishop et al. 2004; Ehlmann et al. 2011a; Viviano et al. 2013). Hence, present-day impingement of these hypothesised aquifers on the surface is hard to rule out and often relies on arguments based around Occam's razor – the simplest hypothesis is the most likely – as proving/disproving the presence of aquifers relies on knowledge of the subsurface that we do not have. So, unlike the other reservoirs of water discussed above, the presence of deep or shallow aquifers (and the cryosphere) remains purely hypothetical and would require significant progress in our understanding of Mars to prove or disprove. Nevertheless liquid water sourced from aquifers is generally perceived to be unlikely to cause present-day or recent (last few millions of years) surface modifications, because of arguments based on structural and topographic context that will be discussed in Section 3.3 (e.g., Treiman 2003). It is generally considered as a viable source of liquid water in the early Amazonian (Baker & Milton 1974; Rodriguez et al. 2015; Dohm et al. 2001).

## 3 What is the evidence for recent liquid water on Mars?

310 As detailed above, our knowledge of the current surface temperature, pressure and humidity conditions on Mars and our expectations derived from Global Climate Models of the recent past, and lack of knowledge regarding the deep subsurface of Mars means that researchers believe that the production of surface liquid water is unlikely. Hence, evidence presented in favour of near-surface liquid water on Mars
is, almost without exception, highly scrutinised. Despite this many surface features have been ascribed
to the action of surface liquid water often in tight analogy with equivalent landforms found on Earth.

Below each of the major landscape elements that have been used to argue for recent surface water on
Mars are described, the evidence supporting its generation by liquid water summarised and the current
assessment of the likelihood of the liquid water hypothesis described. The topics are addressed in a
rough chronological order starting at the present-day and working towards the past.

## 3.1 Slope Streaks and dark dune flows

Three broad types of low albedo downslope-oriented "streaks" resembling a downslope flow have been
observed to form at the present-day on Mars: Recurring Slope Lineae, Slope Streaks and dark dune
flows. All three have been linked to liquid water occurrence at the present-day on Mars. Recurring Slope
Lineae are detailed in the next section.

<figure><img src="image_9.png" /></figure>

Figure 3: Examples of dark dune flows and slope streaks on Mars. Each letter represents a location while each number represents a different time of year. Please note each panel has been individually contrast-stretched to enhance visibility. Panels a1-5 show the evolution of a dark flows emanating from dark spots on the slip face of a linear dune in Richardson Crater (72.0°S, 179.4°E) in the southern hemisphere from MY27 Ls 210-340°. Note the development of the bright halo, the darker albedo at the tip of the flow and the lack of any morphological signature once the frost is removed. Panels b1-5 show the evolution of a

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333 dark flows located on the slip face of a barchan dune at a site informally named "buzzel" (84.0°N, 233.2°E) in the circum-polar erg in the norther hemisphere in MY 29 from Ls 029-087°. Again, note the lack of morphological signature once activity is over. Panels c1-2 show the appearance of new slope streaks on a slope located in the mid-latitude region of Mars (31.1°N, 226.0°E) in MY 31. HiRISE images: a1 PSP_003175_1080 Ls=210.6°, a2 PSP_003386_1080 Ls=220.7°, a3 PSP_003742_1080 Ls=238.1°, a4 PSP_003953_1080 Ls=248.5°, a5 PSP_005931_1080 Ls=340.4°, b1 PSP_007193_2640 Ls=28.8°, b2 PSP_007404_2640 Ls=36.3, b3 PSP_007905_2640 Ls=53.7°, b4 PSP_008248_2640 Ls=65.5°, b5 PSP_008867_2640 Ls=86.6°, c1 ESP_027092_2115 Ls=107° and c2 ESP_031971_2115 Ls=323°. Credit NASA/JPL/UofA.

342 Slope Streaks (Fig. 3c) can be metres to kilometres in total length and they originate on steep slopes 343 (e.g., Sullivan et al. 2001) in the equatorial regions of Mars. They often originate at a point, fan out to 344 reach a maximum width, then terminate in a series of digitate lobes (e.g., Schorghofer et al. 2007) yet 345 their detailed morphology seems to depend on the shape of the underlying slope (e.g., Dundas 2020). 346 Unlike RSL and dark dune flows, Slope Streaks show no incremental growth between images, but simply 347 appear in any given image at their full extent (e.g., Schorghofer et al. 2007). Also, older streaks can 348 appear bright compared to the surrounding terrain (e.g., Aharonson et al. 2003; Baratoux et al. 2006), 349 which is unlike RSL and dark dune flows. Slope Streaks seem to remove a thin layer of material along 350 their path and may push some material aside (Chuang et al. 2007; Phillips et al. 2007). Slope Streaks are 351 found in dusty, low thermal inertia areas of Mars, so are generally assumed to be dry avalanches of dust 352 (e.g. Dundas 2020 sullivan 2001). The involvement of liquid water has been hypothesised, based on their 353 occurrence on slopes whose temperatures can exceed 273K (Schorghofer et al. 2002), a bias in their 354 occurrence towards hotter times of year (Heyer et al. 2019), interaction with topographic obstacles 355 (Brusnikin et al. 2016; Kreslavsky & Head 2009), extension onto low slopes and correspondence with 356 terrains high in Cl and Fe (i.e. availability of salts to produce brines; Bhardwaj et al. 2017). None of these 357 observations are completely diagnostic and dry dust avalanches could have a rheology that also satisfies 358 these constraints (Dundas 2020; Miyamoto 2004; Bhardwaj et al. 2019a).

359 Dark dune flows (Fig. 3a,b) are a strictly seasonal phenomena strongly linked to the defrosting phase 360 during the retreat the seasonal ices in spring and as the name suggests are only found on dark sand 361 dunes (near the poles in both hemispheres). They are metres to hundreds of metres in overall length and 362 develop from dark spots that appear on the seasonal ice cover. They grow incrementally as the seasonal 363 progresses and are no longer visible once the seasonal ice cover has sublimated away, which contrasts 364 with RSL which tend to gradually fade (e.g., McEwen et al. 2011) and slope streaks that can persist for 365 decades (e.g., Schorghofer et al. 2007). Their development follows the same sequence of events: a) 366 development of a dark spot, b) propagation of the flow away from the spot, c) development of bright 367 halos in some cases and d) as propagation progresses the tip of the flow is darker than the source spot. 368 The general consensus is that the initial dark spots are formed by the "Kieffer model" (Kieffer et al. 2006) 369 – this model describes the formation jets of CO₂ gas generated by basal sublimation of the seasonal 370 translucent CO₂ slab ice. The basal sublimation builds pressure under the slab ice that when released by 371 cracking entrains dust/sand that are then deposited around the crack forming a dark spot/fan at the 372 surface. It is proposed that salt can also be included within these deposits and that the dark flows can be 373 explained by the propagation of brines generated as interstitial water from cold-trapped water ice and 374 preferential heating of the sediments forming the dark spot (Kereszturi & Appéré 2014; Kereszturi & 375 Rivera-Valentin 2012; Kereszturi & Rivera-Valentin 2016; Kereszturi et al. 2009; Kereszturi et al. 2011b;
376 Kereszturi et al. 2010; Kereszturi et al. 2011a; Möhlmann 2010; Möhlmann 2004; Möhlmann 2008;
377 Möhlmann & Kereszturi 2010). However, dry mass wasting phenomena driven by sublimation dynamics
378 cannot be ruled out given the morphological, temperature and spectral constraints (Horvath et al. 2009;
379 Hansen et al. 2013; Gardin et al. 2010).

380 Although water has been proposed for the formation of Slope Streaks and dark dune flows, the evidence
381 presented to date favours dry avalanches and CO2 ice sublimation mechanisms respectively.

382 3.2 Recurring Slope Lineae (RSL)

383 Recurring slope lineae are a seasonal phenomenon characterised by downslope propagating relatively
384 low albedo streaks up to hundreds of metres in length and generally metres to tens of metres in width
385 (Fig. 4). They were first identified thanks to repeated HiRISE images and are distinct from slope streaks
386 and dark dune flows by growing at the warmest times of year (for a given slope and slope-orientation)
387 and then fading until the pattern is repeated the following year (McEwen et al. 2011). It is this
388 incremental downhill growth that led researchers to initially interpret these features as being driven by
389 liquid water. This is because slope streaks and wind streaks, which are thought to be produced by dry
390 granular flows (e.g., Chuang et al. 2007; Sullivan et al. 2001) or the removal of dust (Vincendon et al.
391 2015), respectively, do not grow, but are emplaced instantaneously (at the temporal scale of orbital
392 observations). Although there may be mechanisms by which dry granular flows could mimic this
393 behaviour, see Chapter DUNDAS (pages xxx-xxx). RSL occur in exclusively areas that have a low
394 abundance of surficial dust (unlike Slope Streaks which only occur on high albedo dusty slopes) and tend
395 to originate at rocky outcrops towards the top of steep slopes. They are found in dense concentration in
396 the Valles Marineris (VM; Chojnacki et al. 2016; McEwen et al. 2014; Stillman et al. 2017), but also on
397 steep slopes in the northern (Stillman et al. 2016) and southern mid-latitudes (Stillman et al. 2014;
398 Stillman & Grimm 2018) (Fig. 2b) and at tropical latitudes outside of VM (Stillman & Grimm 2018). They
399 often occur within the interior slopes of, on the same slopes as, or on neighbouring slopes to, mid-
400 latitude and equatorial gullies, which are described further in Section 3.3.

401 RSL can be found individually or occur as dense concentrations and they tend to originate at the same
402 place every year. Their overall length varies from year to year and neighbouring RSL can have different
403 growth rates, and in certain locations an RSL can lengthen at the same time its neighbour is fading
404 (Stillman et al. 2017; Stillman et al. 2020). Their path downslope deviates around obstacles (Fig. 4a-c, g-i)
405 suggesting that they have very little momentum and that the flow is thin compared to the obstacles
406 visible at HiRISE resolution (~25 cm/pix). RSL can come together downslope (Fig. 4c, f), split (Fig. 4c, k)
407 and have some limited sinuosity (Fig. 4b, i). In Valles Marineris RSL are spatially and temporally
408 associated with slumps in the loose materials over which they propagate (Chojnacki et al. 2016; Ojha et
409 al. 2017) (Fig. 5). Such slumps extend much further downslope than RSL, have similar albedos, fading
410 timescale, and slopes (27.5° ± 2.7°; Stillman et al. 2020). RSL are most easily identified on materials that
411 are generally smooth at the HiRISE scale (no grains >75 cm), but can be found on slopes covered in clasts
412 (Fig. 4d-f). They can propagate down pre-existing channels (Ojha et al. 2014), including gully channels.
413 They tend to be restricted to slopes above the angle of repose (Dundas et al. 2017), but are found to
414 extend to lower slopes (Schaefer et al. 2019; Stillman et al. 2020; Tebolt et al. 2020). However, RSL form
415 on the upper reaches of scree slopes likely formed via mass wasting and are naturally at the angle of
416 repose. Thus, just because RSL form on such a feature does not require that RSL be interpreted as a dry
417 sediment flow (Tebolt et al. 2020). Isolated dark spots have been observed in rare cases immediately
418 downslope of actively growing RSL, argued to be discontinuous RSL (Fig. 4e) (Ojha et al. 2015; Tebolt et
419 al. 2020). Stillman et al. (2017) interpreted such features as the “porpoising” of water flow, where water

420 flowing under the surface was only sucked to the surface via capillarity in locations of low albedo.
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Figure 4: Examples of Recurring Slope Lineae (RSL), north is up in all panels. (a) RSL in the alcoves of gullies in Corozal Crater in the southern mid-latitudes with (b) showing multiple RSL coming together while deviating around obstacles and (c) showing their termini which often diverge. (a-c) False-colour HiRISE image ESP_023218_1410 where colours have been stretched to best enhance the contrast. (d) RSL extending from rocky substrates onto sand on a ridge in Coprates Chasma in Valles Marineris at the equator with (e) showing their discontinuous downslope extensions and (f) their lower contrast on the upslope rocky materials. (d-f) False-colour HiRISE image ESP_043085_1670 where colours have been stretched to best enhance the contrast. (g) RSL on the interior light-toned layered deposits in Juventae Chasma in Valles Marineris at the equator with (h) showing their propagation over the bedrock with braided sections and (i) showing their extension onto the debris fans, again showing sinuosity and braiding. (g-i) HiRISE image PSP_006915_1760 where stretch has been chosen to best enhance the contrast. (j) RSL on the central peak of Hale Crater in the southern mid-latitudes in HiRISE image ESP_047172_1440 where (k) shows their extension onto lighter albedo "fans". (l) HiRISE image ESP_046605_1440 showing the same location as in panel k, but at a season when the RSL have faded. Image credits: NASA/JPL/UofA.

Concurrent observations of dust devil tracks, boulder tracks, and RSL suggest that the distinctive dark albedo of RSL is a result of the removal of dust (Schaefer et al. 2019) and that annual atmospheric dust deposition allows their reoccurrence the following year. The dust storm of Mars Year 28 and 34 resulted in many more candidate RSL being apparent and for RSL to be in general a lot longer than in previous years (McEwen et al. 2019; Stillman et al. 2014). An alternate hypothesis is that wicking of moisture to the surface or water absorbed into salts could be the cause of the darker albedo (Heinz et al. 2016; Massé et al. 2014). In any case it is agreed that RSL do not engender any significant downslope sediment transport at the HiRISE-scale. As no sediment transport is involved, they are generally thought to represent either subsurface percolation of water, or thin granular flows.

Apart from the various dry hypotheses for RSL formation, discussed further in Chapter DUNDAS (pages xxx-xxx), water-based hypotheses require a mechanism to replenish the supply towards the top of the slope. They therefore focus around two potential mechanisms: i) atmospherically recharged source and ii) origin from aquifer seeps, which are discussed further below.

Mass balance arguments suggest that water ice frosts from the atmosphere are unlikely to be able to recharge RSL source areas, with RSL probably requiring metres cubed of water per metre of headwall, exceeding that which can be supplied from the dry martian atmosphere (Grimm et al. 2014). In addition, the kinetic barrier to melting, detailed in Section 1, would be particularly severe at the relatively warm locations of RSL occurrence in Valles Marineris – meaning the ices would need to somehow be protected by a lag in order to melt. Further, the melting of shallow subsurface ice is difficult as temperatures decrease with depth. If salt were present then melting could occur, but such melting would then flush these salts. This is unlikely because such a shallow subsurface ice would need a mechanism to recharge salt and significant amount of ice so that it could provide a source each year giving the recurrent behaviour of RSL (Stillman et al. 2014).

The surface temperature at which RSL start to grow argues that the water should be briny (i.e., have a depressed freezing point Grimm et al. 2014; Stillman et al. 2016) even though the near infra-red spectral detection of hydrated perchlorates associated with RSL from orbit (CRISM; Ojha et al. 2015) is now generally acknowledged to be due to instrument error (Leask et al. 2018; Vincendon et al. 2019) and no
systematic detection has been made of chlorine bearing salts in thermal spectra (Mitchell & Christensen 2016). The predicted abundance of salts on the martian surface is expected to be high, including chlorine-bearing species (Clark 1981; Gellert 2004; Keller et al. 2007; Kounaves et al. 2010; McSween & Keil 2000) and notably perchlorates (Glavin et al. 2013b; Glavin et al. 2013b; Hecht et al. 2009; Sutter et al. 2017; Toner et al. 2014a) which have remarkably large freezing point depression (e.g., Chevrier & Rivera-Valentin 2012). This abundance has led researchers to propose that RSL could be produced by deliquescence of these salts, i.e. the adsorption of water vapour (from the atmosphere or sublimated from ground ice) leading to the fluidisation of the salt (e.g., Bhardwaj et al. 2019b; Dickson et al. 2013). The relative humidity required for deliquescing Mars-relevant salts is on the order of >13% at 273 K (Gough et al. 2011; Heinz et al. 2016; Nuding et al. 2015; Primm et al. 2017), which for the same reasons as argued above seems high for the dry martian atmosphere. However, once these brines have formed they can form supersaturated and/or supercooled solutions with significant hysteresis prolonging their persistence in the environment (Gough et al. 2016; Primm et al. 2019; Toner et al. 2014b). Although temperatures are highest in the middle of the martian day, the relative humidity is at its lowest, so the early morning and late evening are more conducive to brine-formation by deliquesce (Gough et al. 2011; Primm et al. 2019). If RSL were caused by deliquescent salts, they should have different albedos at different times of day. This behaviour has not yet been observed with when comparing one morning CaSSIS image with an afternoon HiRISE image (Munaretto et al. 2020), but much more extensive analysis is required to substantiate this finding. Additionally, modelled water budgets using meso-scale atmospheric models coupled with deliquescence onto calcium perchlorate salts indicate an upper limit of 1 micron per sol, which is likely much too low to support any wet-dominated RSL formation mechanism (Leung et al. 2020).

Recharging sufficient salts into the source areas could be envisaged via atmospheric deposition, of notably perchlorate salts (Catling et al. 2010; Smith et al. 2014). If the brine comes from underground, there is no recharge problem for the salts, because the aquifer itself would be briny. However, repeating brine flows would be expected to continually deposit salts towards the ends of the RSL and within the subsurface. As no build-up has been detected from orbit there must be some unknown removal mechanism, if this hypothesis is true. The difficulty in recharging sufficient water (Stillman et al. 2016) and the proximity of RSL to fractures observed in the bedrock at which they originate (Abotalib & Heggy 2019) have been used as arguments to support their origin via groundwater. As discussed in Section 3.2 for gullies, a groundwater origin is hard to substantiate with confidence as we have so little knowledge on the structure of the subsurface of Mars (e.g., heatflow, permeability, lithology, fractures, cryosphere state/structure). Many of the most extensive RSL sites are found in locations that are advantageous for groundwater such as: the numerous RSL found within craters in the regional depression of Chyrse and Acidalia Planitia (Stillman et al. 2016), Garni crater on the floor of Melas Chasma (Stillman et al. 2017), and Palikir crater on the floor of the much larger Newton crater (Stillman et al. 2014). These sites could possess unconfined or confined/pressurized aquifers that could hydraulically connect to the surface via low permeability fractures. It should be noted that RSL are also found in places disadvantageous to groundwater access (Chojnacki et al. 2016). Nevertheless, such aquifers would need to possess a freezing depression that is greater than or equal to the mean annual surface temperature so that they could discharge to the surface only during the warmest part of the year (Stillman et al. 2016). Once the brine is discharged, it will then percolate downhill through the regolith. The regolith also must have an impermeable barrier to prevent the brine from seeping back into the subsurface before it shows a surface expression (Mellon & Phillips 2001). Due to the large diurnal temperature variations, the brine is
508 likely to completely or partially freeze at night, thus slowing its advance downhill. As discussed in 509 Chapter DUNDAS (pages xxx-xxx) there are also RSL locations that are not likely to be connected to an 510 aquifer such as the numerous RSL sites found on Coprates Montes (Chojnacki et al. 2016).

<figure><img src="image_22.png" /><figcaption>511
512 Figure 5: 3D rendering of a slump on the same slope as Recurring Slope Lineae (RSL) in Garni Crater. (a) Orthorectified HiRISE image ESP_034672_1685 showing the slope with RSL prior to the slump in MY32, Ls 65°. (a) Orthorectified HiRISE image ESP_035028_1685 showing the slope with RSL and the slump. The image was taken in MY32, Ls 77°, i.e. 26 sols after the image in panel a. Small black arrows point to the location of prominent RSL that grow between the two images. Image credits: NASA/JPL/UofA.
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517 It is possible that liquid water is only the trigger for a motion that is then more granular, for example: the 518 volume or weight change brought about by deliquescence could trigger a granular avalanche (Wang et 519 al. 2019), or humidification-desiccation (water absorption-desorption) cycles could change the cohesive 520 properties of the slope materials (Shoji et al. 2019). Experimental work has shown that boiling water or 521 brines can cause grainflows, or eject pellets of damp material (Herny et al. 2019; Massé et al. 2016; 522 Raack et al. 2017). These more “exotic” processes rely on the contact of liquid water with relatively 523 warm substrates (>280K), as might be envisaged to happen at the edge of a propagating RSL flow-front, 524 and vigorous boiling produces gas in the sediment pore-space capable of fluidising the granular bed, 525 ejecting particles and even mobilising parts of the saturated mass. Such hybrid mechanisms would allow 526 “dry” RSL to follow the seasonal patterns observed for RSL. However, similarly the purely dry 527 mechanisms (Dundas et al. 2017; Schaefer et al. 2019; Schmidt et al. 2017; Vincendon et al. 2019), these 528 granular flows would struggle to imitate the detailed growth histories of RSL (Grimm et al. 2014; Huber 529 et al. 2020) also see Chapter DUNDAS (pages xxx-xxx).

530 In summary, the role of liquid water in Recurring Slope Lineae on Mars is supported by their temporal 531 and spatial patterns of growth at the intra- and inter-seasonal scale, as well as their extension onto 532 slopes lower than the angle of repose. The source of this water is problematic, either from the 533 atmosphere or the subsurface. For brine-flows, a mechanism would be needed to remove the deposited 534 salts. This outstanding problem for the “wet” hypothesis combined with new theories for the “dry” 535 hypothesis has shifted the community’s focus to resolving the issues regarding triggering and growth for 536 the possible “dry” processes, see for example, Chapter DUNDAS (pages xxx-xxx). To date, no advanced 537 modeling has been able to correlate the dry mechanism hypotheses with the observations.
## 3.3 Gullies

The term “gully” for Mars encompass a huge variety of landforms, whose uniting attribute is the presence of a channel through which material is transported from the source alcove to the debris apron or fan (see review by Conway et al. 2019b) (Fig. 6). Gullies are not to be confused with “spur and gully” morphology which was originally reported on Mars in Valles Marineris (e.g., Lucchitta 1987), but in this case the erosional and depositional parts are not connected by a channel. Spur and gully morphology with no channelisation is generally attributed to dry mass wasting processes which build talus cones in oversteepened landscapes (e.g., Thapa et al. 2017) and will not be discussed further here.

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Figure 6: Examples of gullies on Mars. (a) Gullies incised into icy mantling materials on a crater wall in the northern hemisphere. Note the well-formed tributary networks, the abandoned alcoves and multiple generations of fan-deposits. HiRISE image ESP_018895_2410. (b) Gullies in Galap Crater in the southern hemisphere, where the alcoves are incised directly into the bedrock at the crater wall and no icy-mantling materials are present. HiRISE image PSP_003939_1420. (c) “Linear” gullies on the megadune located on the floor of Russell Crater in the southern hemisphere. HiRISE image ESP_038335_1255. Image credits: NASA/JPL/UofA.

Gullies, from the top of the alcove to base of the fan, are typically a kilometre in length (ranging from a hundred metres to several kilometres) and tens to hundreds of metres wide (ranging from decametres to ~kilometre). They were initially interpreted to be a result of liquid water flowing at the surface of Mars, because of their similarity to gullies carved by liquid-water on Earth and because of their latitude distribution and orientation preference (Malin & Edgett 2000). Gullies are found primarily on steep slopes at latitudes polewards of 30° north and south (Balme et al. 2006; Bridges & Lackner 2006; Dickson et al. 2007; Harrison et al. 2015; Heldmann et al. 2007; Heldmann & Mellon 2004; Kneissl et al. 2010) (Fig. 2b) and their variable spatial density in this latitude range is mainly a consequence of the availability of steep slopes (Conway et al. 2017). The sloping terrains where gullies are found comprise the relief-forming structures on Mars at kilometre-scales: impact craters (outer and inner walls, terraces, central peaks/pits), valleys and graben, mesas, buttes, collapse pits and dark sand dunes. At latitudes between 30° and 40° gullies are found primarily on pole-facing slopes and >40° they have a preference for equator-facing slopes, but can be found on slopes with all orientations (e.g., Conway et al. 2017;
Harrison et al. 2015; Heldmann & Mellon 2004). Gullies can be found in the equatorial regions, but they are generally small (hundreds of metres long for metres in channel-width), rare and only reliably identified in HiRISE images (Auld & Dixon 2016; Thomas et al. 2020), hence are not displayed with the larger landforms in Figure 2. It should be noted that Auld and Dixon (2016) classified features with an alcove and an apron but no channel as gullies which deviates from the commonly used Malin and Edgett (2000) definition that requires the presence of a channel. The latitudinal pattern in distribution and orientation of gullies (Fig. 2b) strongly suggests a climatic factor plays a role in conditioning their formation.

<figure><img src="image_24.png" /><figcaption>Figure 7: Examples of changes observed in martian gullies, where the top is the image before the change and at the bottom the image after the change. (a) Dark-toned deposit overlying frost in HiRISE image ESP_027567_1425 (before ESP_022688_1425). (b) Light-toned deposit over an existing fan on HiRISE image ESP_031919_1435 (before ESP_014368_1435). (c) New high-relief deposit with no tonal change in HiRISE image ESP_032078_1420 (before PSP_003939_1420). (d) New channel and deposit in HiRISE image ESP_032011_1425 (before ESP_013115_1420). Note, the before-after pairs are not necessarily the closest in time but are chosen to best illustrate the changes. Image credits: NASA/JPL/UofA.</figcaption></figure>

Worthy of particular attention are "linear gullies" (Fig. 6c) which are a very unusual type of martian gully, lacking a terrestrial equivalent and is only found on steeply sloping dark sand substrates (e.g., Diniega 2021; Diniega et al. 2013; Reiss & Jaumann 2003). "Classic" gullies, as described above, are also found on dunes. Linear gullies comprise a long linear to sinuous parallel-walled groove flanked by levees, which trends downslope and does not change appreciably in width along its length. The source zone is often marked by many smaller tributary grooves and the terminus either blunt or with a series of disconnected
pits, but with no obvious depositional fan/apron. They are often found near to other "classic" gullies and sometimes have morphologies that are transitional with other "classic" gullies (e.g. the channel part resembles a linear gully, whereas the source and deposit apron resembles a classic gully) (e.g., Pasquon et al. 2018). They form a relatively small proportion of the overall gully population, but are one of the most active types of gully at the present-day (Conway et al. 2019b).

Although gullies are too small to be dated directly via crater-size frequency distributions, their recent formation has been inferred by the lack of superposed impact craters. Dating of the terrains that gullies superpose has allowed authors to estimate maximum ages as 1-5 Ma (de Haas et al. 2018a; Reiss et al. 2004; Schon et al. 2009a). Gullies are composite landforms where overlapping deposits demonstrate that gullies have many multiples of formative events (e.g., de Haas et al. 2015a; Johnsson et al. 2014; Levy et al. 2010a). Inverted gully fans/channels and re-incision into mantled alcoves provide further evidence that significant hiatuses occur between different active periods in gullies (e.g., Dickson et al. 2015).

Monitoring of gullies on Mars has revealed that downslope movements of material are ongoing and predominantly occur in local winter (Diniega et al. 2010; Dundas et al. 2019b; Dundas et al. 2015b; Dundas et al. 2012; Jouannic et al. 2018; Malin et al. 2006; McEwen et al. 2007; Pasquon et al. 2019; Pasquon et al. 2018; Pasquon et al. 2016; Raack et al. 2015). Gullies on sand dunes are particularly active and most sites show some kind of motion every year. Mass balance calculations on dune gullies has highlighted that these features can be formed entirely in as little as a few hundreds of years (Pasquon et al. 2018). Only a few other sites show such cadences, with the majority of gullies showing activity having only one change within the last few decades of monitoring (e.g., Dundas et al. 2019b). Movements mostly take the form of relatively high or low albedo deposits on fans or within gully-channels (Fig. 7b), with only some identifiable by their relief (Fig. 7c,d). In a few cases the source of the sediment can be tracked back to an identifiable scarp (e.g., de Haas et al. 2019; Raack et al. 2020). Channels can be evacuated, or back-filled with sediment and new channels carved into fan-surfaces. Sites with gullies also often host RSL (Fig. 4a-c) and often small-scale equatorial gullies are also found at RSL sites (Fig. 4g-i).

Gullies have been proposed to form via both dry and wet processes, and a summary of the arguments in favour of the dry hypotheses can be found in Chapter DUNDAS (pages xxx-xxx). Due to present-day activity occurring during local winter it is very unlikely that these motions are triggered by liquid water. The only exception is that the recent activity in linear gullies does occur at a time of year when water ice can be cold-trapped and surface temperatures are rapidly rising towards the melting point (Pasquon et al. 2016), but even in this case the link with CO2 ices provides a simpler explanation and is more widely accepted (see Chapter DUNDAS, pages xxx-xxx). It should be mentioned that some present-day activity characterised by the bright deposits, occur at times and places where CO2 ices have not been detected, but H2O ices have (Vincendon 2015), suggesting that water ice might be playing a role, although such deposits are found to be consistent with dry granular flows (Kolb et al. 2010; Pelletier et al. 2008). Whether the sediment transport observed at the present-day is representative of the processes that have formed the entire gully-landform still remains open to debate as gullies have a number of characteristics that are consistent with formation by flowing liquid water, as we detail below.

Martian gullies share many morphological characteristics with water-driven steepland erosional systems on Earth. Many of the detailed elements of martian gully morphology can be identified within larger terrestrial catchments, but also erosion-deposition systems that share a similar morphology can be found in permafrost to arid regions on Earth (Costard et al. 2007; Dickson et al. 2018; Hartmann et al.
631 2003; Hauber et al. 2018; Heldmann et al. 2010; Sinha et al. 2018). Sediment transport through such
632 systems is typically via a cascade of different processes, including mass wasting or landsliding, debris
633 flow and fluvial transport (e.g., Cavalli et al. 2013; Heckmann & Schwanghart 2013; May & Gresswell
634 2004). The upper part of martian gullies is often tributary in nature, forming an organised, hierarchical
635 network, a typical organisation of fluvial systems on Earth. The lower part of martian gullies often
636 comprises a fan, where multiple deposits overlap, a characteristic shared by alluvial fans on Earth (de
637 Haas et al. 2016; de Haas et al. 2018b). Fresh individual deposits can be characterised by lobate margins
638 and levees flanking the channel (Johnsson et al. 2014; Lanza et al. 2010; Levy et al. 2010a; Sinha et al.
639 2020), which are characteristic of debris flow processes on fans on Earth (e.g., Blair & McPherson 2009;
640 de Haas et al. 2015c). In-channel deposits can show braiding, streamlining, sinuosity, abandoned channel
641 segments, terraces, all typical features of sediment-charged streams on Earth (e.g., Bakker et al. 2019).
642 Topographically, both in terms of along-channel profiles and in terms of contributory-characteristics
643 martian gullies are indistinguishable from terrestrial systems dominated by debris flow or fluvial
644 processes (Conway et al. 2015; Conway et al. 2011b; Conway & Balme 2016; Hobbs et al. 2017; Yue et al.
645 2014). Stratigraphic analysis has revealed that the deposits laid down in martian gully fans indicate a
646 dominance of debris flow processes, even though typical debris flow surface morphologies, such as lobes
647 and levees, are absent, probably as a result of surface erosion by wind (de Haas et al. 2015b).

648 Numerical models and experimental simulations both conclude that once liquid water is generated, it
649 would be sufficiently stable to flow the whole length of the gullies and to transport sediment (Bargery &
650 Gilbert 2008; Conway et al. 2011a; Heldmann et al. 2005; Herny et al. 2019; Parsons & Nimmo 2010;
651 Raack et al. 2017). Its metastable state may result in some transport mechanisms that are rarely active
652 on Earth, particularly those as a result of boiling, which can result in enhanced sediment transport for
653 relatively limited quantities of water (Herny et al. 2019; Massé et al. 2016; Raack et al. 2017).

654 No matter the source for the liquid water, the proposed release mechanism is a result of the insolation
655 conditions experienced by Mars under moderate to high obliquity conditions (>35°), when peak
656 temperatures are experienced by steep pole-facing slopes in the mid-latitudes (Christensen 2003;
657 Costard et al. 2002; Williams et al. 2009; Williams et al. 2008). Three sources of water have been
658 considered as possible: i) groundwater, ii) ground ice, or iii) surface ice precipitates (snow, or frost).

659 *Groundwater*: Both deep (Gaidos 2001; Grasby et al. 2014) and shallow sources (Heldmann & Mellon
660 2004; Marquez et al. 2005; Mellon & Phillips 2001) have been proposed as the source for gullies, where
661 in both scenarios seeps of water/brine occur when the incoming radiation is sufficient to melt the plug of
662 ice formed where the aquifer reaches the surface, and falling winter temperatures reseal this plug.
663 Where the radiation is too strong, water will evaporate rather than flow and where the radiation is not
664 strong enough, the plug will never melt – these factors explain the latitude orientation trends shown by
665 gullies (Fig. 2c). Both sources struggle to explain the emergence of gully alcoves at the top of isolated
666 topography, where shallow aquifers would not be sufficiently voluminous to repeatedly supply sufficient
667 water and deep aquifers would be improbable to emerge (e.g., Balme et al. 2006; Conway et al. 2019b;
668 Treiman 2003). Although, recent arguments derived from RSL-research (Section 3.1) maintain that deep,
669 pressurised aquifers could emerge at these elevated locations if a permeable pathway exists (Stillman et
670 al. 2016).

671 *Ground ice*: Redistribution of the water ice from the polar caps to lower latitudes under conditions of
672 high orbital obliquity is thought to explain the prevalence of ground ice down to 50° north and south
673 (Jakosky & Carr 1985; Madeleine et al. 2014; Mellon & Jakosky 1995) and its discontinuous presence to
674 30°N and S (Kreslavsky & Head 2000; Milliken et al. 2003; Mustard et al. 2001). The same high obliquity
675 conditions could be conducive to its melting where peak temperatures are experienced in the subsurface
676 on pole-facing slopes in the mid-latitudes (Costard et al. 2002; Kreslavsky & Head 2003). In this scenario
677 the distribution and orientation of gullies with latitude is a factor of both the availability of ground ice
678 and the conditions required for its melting. Mass balance arguments show that ground ice is being lost
679 during gully-formation (Conway & Balme 2014), but gullies are found in zones without obvious presence
680 of ground ice (de Haas et al. 2018a; Johnsson et al. 2014), which argues against this hypothesis.
681 Conversely, the abrupt absence of equator-facing gullies at 40° north and south matches surprisingly well
682 with the anticipated absence of ground ice at these locations (Conway et al. 2018b).

683 *Snow/frost*: The formation scenario for snow and/or frost (de Haas et al. 2015a; Kossacki & Markiewicz
684 2004; Williams et al. 2008) is similar to that for ground ice, but the source is seasonal surface deposits,
685 rather than ice stored within the soil. For one particular site the debris flow cadence revealed by the
686 deposit volumes and the age of the host crater is consistent with millimetres of snow over the
687 catchment, which is broadly in accord with climate model predictions for snowfall at high obliquity (de
688 Haas et al. 2015a). Whether atmospheric conditions in the past were conducive to melting at the surface,
689 however, is a matter of debate.

690 In summary, although gullies on Mars strongly resemble terrestrial gullies carved by liquid water, it
691 remains under debate as to whether the same processes were active on Mars. There is growing evidence
692 that present-day activity almost certainly does not involve liquid water, but whether present-day
693 processes are representative of the processes that have contributed to the construction of the whole
694 gully-landform is still debated. If current modification processes of gullies cannot explain their entire
695 formation, liquid water at high obliquity is the next most likely candidate to explain the formation of
696 martian gullies. Periglacial landforms

697 Many landforms tightly associated with periglacial conditions on Earth have also been reported on Mars,
698 including: solifluction lobes, pingos, low-centred polygons, ice-loss depressions and sorted patterned
699 ground. For a in-depth review of these features on Mars, please refer to Balme et al. (2013). Periglacial
700 conditions are those which are dominated by the action of freeze-thaw processes and do not necessarily
701 imply permafrost or glacial conditions (French 2013). Integral to periglacial conditions is the occurrence
702 of freeze-thaw cycles and the formation of an active layer – that is a layer on top of the permafrost that
703 undergoes annual thaw – hence in apparent contraction with our understanding of the stability of liquid
704 water on Mars. Implicit for the formation of periglacial landforms is the presence of ground ice, which as
705 described in Section 2.2, which is a pre-requisite relatively easy to fulfil on Mars.

706 Below, we summarise the observations pertaining to each of these features and the way in which liquid
707 water is thought to be involved in their generation and then introduce the concept of a landscape
708 assemblage.

## 3.4 "Solifluction" lobes

710 Lobate forms found on hillslopes on Mars have a strong resemblance in form and scale to solifluction
711 lobes on Earth (Gallagher et al. 2011; Johnsson et al. 2018; Johnsson et al. 2012; Soare et al. 2018b) (Fig.
712 8). Like their terrestrial analogues they can be sheet-like or tongue-like, extending in the downslope
713 direction and are bounded at their lower end by a scarp or "tread". Their margins can be picked out by
714 relief or by clasts. The martian lobes can be larger in both width and length than solifluction lobes on
Earth (generally up hundreds of metres), but their ranges of sizes do overlap and are typically tens of metres wide and long (Gastineau et al. 2020). Even though individual lobes are small, they occur in groups that can span hundreds of metres to kilometres across and downslope. Martian lobes have the peculiarity that they are found on slopes above 10° and on average ~25°, whereas terrestrial solifluction lobes can be found on all slopes up to 35° (Gastineau et al. 2020). They occur at latitudes poleward of 60° in the northern hemisphere and at latitudes of poleward of 40° in the southern hemisphere (Fig. 2c) and are less common in the southern hemisphere (Johnsson et al. 2018).

<figure><img src="image_25.png" /><figcaption>Figure 8: Examples of lobate forms and pingo-like mounds. (a) Lobate forms outlined by clasts on the wall of Heimdal Crater (Gallagher et al. 2011). Downslope is down-image. HiRISE image PSP_009580_2485. (b) Lobate forms marked by relief and overprinting polygonal terrain on an impact crater wall in the northern hemisphere (Johnsson et al. 2012). Downslope is down-image. HiRISE image PSP_010077_2520. (c) Pingo-like mound near Moreux Crater (see Chapter SOARE, pages xxx-xxx) with prominent summit cracks. HiRISE image ESP_058140_2225. Credits: NASA/JPL/UofA.</figcaption></figure>

Solifluction lobes on Earth are sheets of material that move as a result of differential downslope movement caused by freeze-thaw cycling in the active layer (Matsuoka 2001). A range of processes lead to the formation of solifluction lobes on Earth, including: gelifluction, plug flow, frost creep and retrograde movement (Matsuoka 2001). The influence of each of these processes on the morphology of the lobes has not been isolated, but each requires extensive thaw and considering martian lobes cover hundreds of metres of areal extent, this would require equivalently extensive thaw to prevail. The age of these features is uncertain, but is presumed to be late Amazonian as they are not overlain by any other feature and there is some evidence of motion today in the form of boulder motions on the same slopes (Dundas et al. 2019a). Alternatives to thaw for generating lobate forms have not yet yielded convincing results (Gastineau et al. 2020), but some kind of creep triggered by CO₂ ice-related processes has been suggested, but remains to be substantiated (Dundas et al. 2019a).

## 3.5 Pingo-like mounds

Mounds with spatial extents of tens to hundreds of metres are commonplace on Mars and have been attributed to the action of a range of processes, including differential erosion, volcanism (magmatic or sedimentary), inverted impact craters and pingos (e.g., Bruno et al. 2006; Burr et al. 2009a). Of particular interest for this discussion are those that have been identified as pingos (Fig. 7) (Balme & Gallagher 2009; Burr et al. 2009b; Dundas et al. 2008; Dundas & McEwen 2010; Soare et al. 2013b; Soare et al. 2013a; Soare et al. 2014a; Soare et al. 2005), see also Chapter SOARE (pages xxx-xxx). These landforms
747 have been reported mainly in the mid-to-high-latitudes on Mars with the exception of one report near
748 the equator (Balme & Gallagher 2009).

749 Pingos on Earth can be formed by the impingement of permafrost upon a near-surface liquid water body
750 (talik) generating hydrostatic pressure (closed system) or by artesian pressure developed in a talik
751 connected to the surface (open system). Typically closed system pingos are developed when an alas, or
752 thermokarst lake, catastrophically drains and the permafrost re-advances into the saturated sediments
753 (e.g., Mackay 2002). Open system pingos develop where sub-permafrost meltwater is pushed to the
754 surface by an elevated hydrostatic gradient generated by surrounding topography, for example in the
755 forefields of retreating glaciers (Christiansen 1995; French 2013). In both cases massive ice is formed in
756 the core of the pingo which pushes up the surface above it, forming a mound. The mounds can be
757 circular, elongate or form complexes in planform. On Earth, pingos can grow to heights of tens of metres
758 with widths of hundreds of metres to nearly a kilometre (Gosse & Jones 2011; Jones et al. 2012). As they
759 grow via injection of water and deformation of the ground, they often display summit cracks where the
760 deformation is greatest (e.g., Mackay 1987). Their flanks slopes can increase beyond the angle of repose
761 then causing mass wasting of the flanks. The cracks in the overburden can engender the destabilisation
762 of the interior ice, forming a summit pit. If the ice at the core melts then the pingo can completely
763 deflate leaving a shallow raised rim depression or pingo scar (Flemal 1976).

764 Similarly, pingos on as Mars have been differentiated from other mounds, such as rootless cones,
765 principally by their summit cracks (Fig. 8) and pits, with some displaying flank mass wasting and irregular
766 planforms. Such mounds have been reported in closed topographic environments, such as crater floors
767 (Soare et al. 2013a), as well as at the base of slopes in more open topographic environments (Soare et al.
768 2014a). The implication of pingos on Mars is that there was sufficient melt to generate the talik, but no
769 long-lived sub-surface water is required once they are formed. The age of the mounds on Mars is
770 unknown, but they are not superposed by any other feature, so are presumably Late Amazonian in age.
771 No alternative mechanism for ground-uplifting has been proposed as an alternative to the liquid water
772 hypothesis, but the interpretation of these mounds as being a result of heave (and not one of the other
773 potential and numerous mound-forming processes) remains under question, as liquid water is
774 considered unlikely.

## 3.6 Patterned ground

775 Patterned grounds on Earth are characterised by the separation of the surface soils into relatively fine
776 and clastic domains with regular, repeating patterns. These take the form of, sorted stone circles,
777 labyrinths, polygon-nets or piles, which can evolve into clastic stripes on sloping terrain. The ratio
778 between the fines and clasts, the amount of confinement and the slope are all factors that play a role in
779 determining the patterns that are expressed at the surface (Kessler & Werner 2003). The mechanism
780 that leads to the sorting is a feedback that involves frost heave combined with differing thermal inertias
781 of saturated soil and stones (Hallett 1990) and necessarily involves freeze-thaw cycles. On Mars surface
782 patterns resembling those of sorted patterned ground on Earth have been identified at the equator
783 (Balme & Gallagher 2009), but more pervasively in the northern plains at >35 °N (Barrett et al. 2018;
784 Barrett et al. 2017; Gallagher et al. 2011; Gallagher & Balme 2011; Soare et al. 2018a), but also in the
785 southern hemisphere (Soare et al. 2016). The martian landforms include clastic polygons, stripes and
786 boulder piles (Fig. 9) and tend to be at the larger end of the scale of the sorted patterned grounds
787 observed on Earth, with length-scales of tens of metres, as opposed to metres as usually seen on Earth
788
789 (Barrett et al. 2018). The patterns follow the expected transition to stripes on sloping terrain on Mars as
790 they do on Earth.

791 The presence of these landforms on Mars suggests a climatic episode where freeze-thaw cycles were
792 common allowing the development of these landforms. An alternate CO2 ice driven process is described
793 in Chapter DUNDAS (pages xxx-xxx) (Orloff et al. 2011; Orloff et al. 2013), but no experimental or field
794 evidence is available to test its efficacy, hence remains difficult to assess its likelihood. Whether these
795 landforms are active at the present-day is unknown and their age is difficult to determine but given their
796 uncratered and pristine appearance it is assumed to be late Amazonian.

<figure><img src="image_26.png" /><figcaption>Figure 9: Examples of sorted forms on Mars. (a) Sorted stripes in Heimdal crater (Gallagher et al. 2011), where the stripes are picked out by clasts, downslope is to the bottom of the image. HiRISE image PSP_009580_2485. (b) Sorted circles on the plains outside Heimdal Crater (Gallagher et al. 2011). HiRISE image PSP_009580_2485. (c) Sorted nets near Lomonosov Crater (Barrett et al. 2017), where clasts are arranged into polygonal patterns. HiRISE image ESP_017131_2485. Image credits: NASA/JPL/UofA.</figcaption></figure>

803 3.7 Low-centred polygons

804 In Section 2.3.2 we described the prevalence of polygonal patterns in the martian mid-to-high latitudes,
805 which are believed to be a direct result of thermal contraction cracking of perennially frozen ground
806 (Mellon 1997). The polygons are generally metres to tens of metres in diameter (Mangold 2005). In
807 many areas the centres of these polygons are at lower elevations than their edge often with a double
808 raised rim (Levy et al. 2010b; Levy et al. 2009a; Soare et al. 2014b; Soare et al. 2018b) (Fig. 10d-f), "low-
809 centred polygons", a phenomenon on Earth that is linked to active ice-wedge or sand-wedge polygons.
810 Ice-wedge polygons form when liquid water penetrates cracks formed by thermal contraction and
811 expands to both force the crack further apart and pushes up the material above it (MacKay 2002). Sand-
812 wedges form when sand is repeatedly transported into the contraction crack and raised rims can be
813 formed when the thermal re-expansion of the ground pushes against the additional material that has
814 accumulated in the wedge causing the soil to be displaced upwards at the margin (Murton et al. 2000;
815 Pewe 1959). Without in situ investigation it remains unknown if the low centred polygons on Mars
816 comprise ice or sand wedges and however, the presence of ice-wedges has been inferred by the
817 landscape context of the low-centred polygons, a concept discussed further in Section 3.3.6. On their
818 own low-centred polygons are ambiguous and have not been used to argue for the action of liquid
819 water, but have been incorporated as part of a landscape assemblage that has been used to argue for
820 the action of liquid water, see Section 3.3.6.
3.8 Ice-loss depressions

We refer readers to Chapter VIOLA (pages xxx-xxx) for a complete description of landforms interpreted to be a result of ice-loss collapse processes on Mars, but for completeness we provide an abridged description here as well. These landforms are often referred to as "thermokarst", which on Earth refers to landscapes created by collapse induced by thaw (increase in temperature), and on Mars this is likely not the case, so we do not use this term for Mars. "Cryokarst" is another term that has been used for these landforms (Costard 2007), but is a term usually restricted to collapse features engendered by the loss of glacial ice via melting (e.g., Corbel & Gallo 1970), hence we do not apply this term either. For collapse to occur the ground must contain excess ice, that is ice in excess of the pore-space, hence when it is lost a decrease in volume occurs. On Earth, ice-loss in thermokarst occurs via thaw, but on Mars ice-loss is more likely to occur by sublimation (e.g., Dundas et al. 2015a). This is in accord with observations, as no evidence of lakes or drainage channels have been observed. On Earth, thermokarst landforms are varied (e.g., Kokelj & Jorgenson 2013), but of particular relevance to Mars are three features: terrain comprised of circular to oblate depressions often filled with water "alases" resulting from the degradation of excess ice, high centred polygons resulting from ice-loss at the polygon margins and beaded streams that are a result of the melting of ice in ice-wedge polygons. So called "scalloped depressions" are found on Mars (Fig. 10b,c,e) and parallels have been drawn between these and alases on Earth (e.g., Lefort et al. 2010; Soare et al. 2008; Ulrich et al. 2010). These depressions are found polewards of 40°N (Orgel et al. 2018; Ramsdale et al. 2018; Séjourné et al. 2018) and poleward of 35°S (Voelker et al. 2017; Zanetti et al. 2010). A related feature is expanded craters (Fig. 10a), which are circular resembling impact craters yet have an identifiable terrace (Viola et al. 2015; Viola & McEwen 2018), and they have similar latitudinal distribution to scalloped depressions. High-centred polygons (Fig. 10e) are found both in association with other ice-loss depressions on Mars and as isolated landforms, but as on Earth are associated with ice-loss at the polygon margins (e.g., Levy et al. 2009b; Soare et al. 2014b). On Earth, this loss can be from an ice-wedge (e.g., Ulrich et al. 2011), but also simply general loss of underlying massive or excess ice focused via the thermal contraction crack at the polygon margin (Marchant et al. 2002). Polygon junction pits (Fig. 10b) are observed in Utopia Planitia (Costard et al. 2016; McGill 1986; Soare et al. 2012; Soare et al. 2011) and parallels have been drawn between these expanded polygon margins and degraded ice-wedge polygons on Earth which form beaded-streams. All these ice-loss features are found in the mid-latitudes and also show concentrations at certain longitudes: Utopia and Acidalia Planitae in the northern hemisphere and Malea Planum in the southern hemisphere. The age of the ice-loss depressions is not easy to determine – they superpose all other landscape features, so are thought to be Late Amazonian (see Chapter VIOLA, pages xxx-xxx).

Although it is likely that the ice-loss depressions on Mars do not involve the generation of liquid water, their presence is a strong indicator of excess ice and the generation of excess ice on Earth involves the presence of saturated ground (e.g., Rempel 2007). Excess ice is differentiated from glacial or massive ice as it requires development of ice within the soil column, rather than deposition of ice directly at the surface. Ice-loss terrain cannot develop in glacial ice as there would not be sufficient material left behind after thaw/sublimation to maintain the landform relief. Porous ice (analogous to firn) has been inferred to be present in areas with ice-loss terrain to depths of around a hundred metres from radar sounding data (Bramson et al. 2017; Bramson et al. 2015; Stuurman et al. 2016), and could have up to 20% lithic material, however this is at odds with geomorphic evidence showing that ice could not have been precipitated from the atmosphere (e.g., boulders at the surface; Sizemore et al. 2015). Enriching the martian soil column to tens of metres of depth with ice in the absence of liquid water is challenging
because ice diffused into the pore-space tends to seal the underlying soil column from further deposition (Jakosky 1983; Zent et al. 1986), even if migration through thermal contraction cracks is taken into account (Fisher 2005). Segregation ice has been observed at the Phoenix landing site and can be generated without thaw (Mellon et al. 2009; Sizemore et al. 2015), but such diffusive processes are only thought to be able to produce excess ice in the first metre of the soil.

Note, that a porous medium with salt and ice in its pores will contain some residual liquid. The amount of unfrozen water depends on the temperature, salt concentration, salt type, and surface area of the medium. This water stays unfrozen due to freezing point depression due to the dissolved salts and because of the effects of surface energy along curved interfaces (Rempel 2012; Sizemore et al. 2015). When a porous media with an appreciable amount of unfrozen water and permeability below freezing undergoes temperature cycling near 273 K, it can form segregated ice (i.e. ice lenses). This process is called cryosuction and occurs under a temperature gradient, where unfrozen water is suctioned to colder temperatures when the liquid veins in the colder soil constrict due to the colder temperature. This new water is then frozen and more water is drawn into the colder porous media. With the very cold temperatures of Mars, temperature cycling near 273 K only occurs within the first metre of the soil.

Hence, although ice-loss depressions are not thought to be engendered by melt, it is still an open question as to whether the generation of the terrain hosting these depressions implies enrichment of the ground in ice via migration of unfrozen water. The timing of the emplacement of the terrain hosting the ice-loss depressions is hard to determine, but is most likely to be during the Amazonian (Soare et al. 2015; Willmes et al. 2012) (see Chapter VIOLA, pages xxx-xxx and Chapter SOARE, pages xxx-xxx), although could be older.

<figure><img src="image_27.png" /></figure>
887 Figure 10: Examples of martian ice-loss depressions and different polygonal patterns. (a) Expanded
888 secondary craters, where a polygonised basin can be seen extending away from the central bowl-shaped
889 depression (REF). HiRISE image ESP_028411_2330. (b) "Scalloped" depressions in Utopia Planitia
890 intersected by polygon-junction pits. HiRISE image PSP_007740_2250. (d) "Wavy" and polygonised
891 depressions associated with gullies in the Argyre Basin. HiRISE image ESP_040974_1395. (e) Low-centred
892 polygons on the inter-crater plains in the high southern latitudes. HiRISE image ESP_022381_1100. (e)
893 Scalloped depression in Utopia Planitia with low-centred polygons near its wall and high-centred
894 polygons on its floor. HiRISE image ESP_034164_2260. (f) Polygon margins marked by a double-raised
895 rim. HiRISE image PSP_005821_1095. Image credits: NASA/JPL/UofA.

896 3.9 Landscape Assemblage

897 Above, we have examined each landform independently and examined the likelihood that each could be
898 a result of processes involving liquid water. However, landforms do not occur in isolation and both the
899 context and the association between landforms can favour one interpretation over another. This concept
900 of a landscape assemblage has been used to argue for the involvement of liquid water in the generation
901 of landscapes on Mars (Balme et al. 2013; Gallagher & Balme 2011; Hauber et al. 2011a; Hauber et al.
902 2011b; Soare et al. 2014b; Soare et al. 2018b; Soare et al. 2018a; Soare et al. 2016; Soare et al. 2005)
903 (see Chapter SOARE, pages xxx-xxx), where gullies, sorted patterned ground, lobate forms, low-centred
904 polygons, pingos and/or ice-loss depressions are found in close proximity suggesting a common suite of
905 processes in their origin. Figure 11 outlines such a scenario where freeze-thaw cycles within the
906 materials in the crater wall, brings about sorting (to form stripes) and gelifluction to create solifluction
907 lobes – extreme melting events lead to generation of gullies, whose melt water drives the formation of
908 pingos at the base of the slope. The context of many of these assemblages is interpreted to be post-
909 glacial, because of the following attributes: “washboard” texture on slopes (indicative of downslope
910 deformation), polygonal surface textures and ice-loss-depressions (indicating high ice content),
911 expanded fractures (which could be crevasses), all upslope of arcuate ridges with spatulate depressions
912 (interpreted to be moraines, e.g., Berman et al. 2005) and often stratigraphically above older glacial units
913 (Head et al. 2008; Jawin et al. 2018).

914 The fact that CO2 ice processes and sublimation landscape assemblages are not found on Earth means an
915 equivalent comparison cannot be made, a point that will be discussed further in Section 4.
<figure><img src="image_28.png" /></figure>

Figure 11: Sketch of landscape assemblage on a crater wall, with (from top to bottom), sorted circles leading into sorted stripes then solifluction lobes. Hillslope is incised by gullies, whose fans are partially overprinted by extensional fractures. At the base of the gullies are arcuate ridges and in the intervening basin, a pingo.

## 3.10 Evidence for melting of martian icecaps or glaciers

Evidence for melting of martian icecaps or glaciers spans from the present-day, to contemporaneous with gullies (millions of years ago), contemporaneous with extant martian glaciers (hundreds of millions of years ago) to much earlier in martian history (billions of years ago). Orosei et al. (2018) and Lauro et al. (2020) reported radar evidence which suggested the presence of lakes of liquid brine underneath the south polar cap of Mars at the present-day. Recent modelling work has highlighted that the position of the proposed lake is not located in a local depression, but is connected to the hydrological system, hence is more likely a pocket of brine caused by locally elevated geothermal flux (Arnold et al. 2019). While
MARSIS detects a significant radar return, it is unfortunate that the higher frequency radar sounder SHARAD has not detected this anomaly. SHARAD cannot always image the floor of the SPLD, however, given such a large signal one would expect a detection. Hence, although not substantiated by other evidence, it remains plausible that liquid water can be generated at the base of the martian polar caps, even under present-day conditions. Various lines of evidence already point to the fact that the south polar cap has undergone basal melting in the Early Hesperian – Late Noachian (e.g., Kress & Head 2015). Sinuous ridges emanating away from the boundary of the south polar cap have been interpreted to be eskers – deposits laid down in sub-glacial channels (Butcher et al. 2016; Kress & Head 2015; Scanlon et al. 2018; Scanlon et al. 2015). Scenarios in which elevated geothermal heat flux is required in the form of volcano-ice interaction (Ackiss & Wray 2014; Ghatan & Head III 2002; Wray et al. 2009) and passive climate-driven melting with a generally larger Noachian geothermal heat-flux (Fastook et al. 2012) have been explored and found to be plausible.

<figure><img src="image_29.png" /></figure>
943 Figure 12: Evidence of basal melting of glaciers on Mars. (a) Sinuous ridge in Phlegra Montes (Gallagher & Balme 2015) interpreted as an esker, whose parent glacier is located off the bottom of the image. CTX Images G20_026224_2116 and P22_009583_2132. (b) Sinuous ridge emerging from parent glacier located in lower half of image in Tempe Terra (Butcher et al. 2017). CTX image P05_002907_2258. (c) Crater in the southern mid-latitudes in CTX image G09_021563_1427 showing evidence of glacial erosion, where the glaciated pole-facing slope in (d) has a lower slope-angle with the bedrock spurs planed off (arrows) directly above arcuate ridges interpreted to be moraines, and the non-glaciated equator-facing slope in (e) has preserved bedrock spurs exhibiting a higher slope angle (Conway et al. 2018a). (d-e) HiRISE image PSP_006663_1425. Image credits: NASA/JPL/UofA/MSSS.

952 Sub-glacial or sub-ice-cap melting has also been proposed as an alternative to "the warm wet early Mars" explanation for the valley networks (Section 1) which are a global surface feature (e.g., Cassanelli & Head 2015; Fastook & Head 2015; Weiss & Head 2015; Wordsworth et al. 2013; Grau Galofre et al. 2020). In contrast, the presently visible glaciers on Mars are generally believed to be cold-based, that is, their interior temperature never rises above the freezing point of water, hence their deformation progresses via the plastic deformation of ice integrated over long time periods. As described in Section 2.2, the martian mid-latitudes are host to a huge number of glacial forms (Fig. 2d), which are believed to be several tens to hundreds of million years old (see Hepburn et al. 2019 and references therein). Some evidence for surface melting of these features comes in the identification of tens of small supraglacial and pro-glacial channels (Fassett et al. 2010), but such channels do not imply the long-lived liquid water involved in basal melting. Recent work has questioned the cold-based paradigm for Amazonian glaciers in two ways: firstly examples of Amazonian debris covered glaciers directly associated with eskers have been reported (Butcher et al. 2017; Gallagher & Balme 2015), and evidence has been found for enhanced erosion of mid-latitude crater walls which is consistent with wet-based glaciation in the last 5-10 Ma (Conway et al. 2018a).

967 In Phlegra Montes and Tempe Terra, sinuous ridges emerge from extant debris covered glaciers found within fault-bounded graben (Fig. 12). These ridges are interpreted as being eskers, not just because of their intimate association with the glaciers, but because their morphological and topographic relationships are consistent with expectations from terrestrial glaciers/ice sheets (Butcher et al. 2020; Butcher et al. 2017; Gallagher & Balme 2015). The formation of eskers requires substantial amounts of liquid water and its fate once it emerged from under the ice would have likely been quick evaporation and freezing, although no geomorphic evidence is apparent to support or refute this inference. It should be noted that on Earth catastrophic outbursts from sub-glacial conduits do not produce terminal deposits, in contrast to long-lived more stable flows (Burke et al. 2012; Butcher et al. 2020). As both systems are within fault-bounded graben, elevated geothermal heat flux associated with the crustal-thinning could be the reason that only these isolated individual esker systems have been found.

978 Analysis of high resolution digital elevation models has revealed systematic lowering of the bedrock slopes of pole-facing crater walls in the mid-latitudes intimately associated with the removal of bedrock spurs and transformation of the texture of the exposed bedrock (Conway et al. 2018a)(Fig. 12). The amount of horizontal retreat required to explain the observed slope-lowering ranged up to 10² m/Myr and this elevated value is consistent with erosion rates of wet-based glaciers on Earth. Cold-based glaciers on Earth have erosion rates of tens of metres per million years and rates on Mars are expected to similar or lower because of the lower gravity and lower mean average temperature making the ice more rigid (Karlsson et al. 2015). The parts of the craters affected by the slope-lowering are intimately
associated with the presence of a "pasted-on" mantling deposit located downslope of the removed bedrock, which can be interpreted to be an ice-rich till-like deposit due to the presence of downslope lineaments (striations). Where this mantle-deposit interacts with older glacial deposits (LDA or CCF) located on the crater floor there are arcuate ridges interpreted to be moraines (Arfstrom & Hartmann 2005; Berman et al. 2008; Berman et al. 2005; Head et al. 2008; Jawin et al. 2018; Whalley & Azizi 2003).
On Earth, terminal moraine ridges are formed by thrusting of the basal sediments towards the glacial surface, which is facilitated by the presence of liquid water (e.g., Benn & Evans 2010), but recent research has revealed they can occur due to deceleration in entirely frozen cold-based glaciers (Fitzsimons & Howarth 2020).

The implication of this work is that a limited amount of melting occurred within the top decametres of the ground in the last 5-10 Ma over a widespread area on Mars. The exact trigger for this "event" is unknown, but seems to be related to the shift of Mars to lower average obliquity at 4.5 Ma (Laskar et al. 2004; Laskar et al. 2002). Another possibility is that periglacial processes have caused the observed slope-lowering (which would still involve liquid water) or some other process driven by H₂O or CO₂ frost-driven rock-breakdown, but at present no viable mechanism has been proposed. Unlike the phase-change from liquid water to solid water, direct condensation from gas to solid does not involve a volume change so rock-breakdown could only be engendered by deposition of frost within cracks opened by thermal contraction whose closure on warming is then limited by the presence of the ice. It should be noted that direct condensation of water ice at extremely low vapour pressures has been shown to always pass via the liquid phase and only in porous particles in the terrestrial atmosphere (David et al. 2019). Given the close association of this suite of landforms with gullies it is tempting to implicate melting for the whole suite of landforms, but whether this is reasonable is certainly debatable.

## 3.11 Assemblages of pitted cones, flows, ridges and fractures

Sedimentary volcanism is the mobilisation of sediments by groundwater circulation and is usually engendered on Earth pressurised aquifers caused by dense sediments overlying less-dense saturated sediments. Most eruptions are facilitated by gas production (usually methane) and/or tectonic activity. When this process involves clay-sized sediments it is referred to as mud-volcanism. Pitted cones interpreted to be mud-volcanoes have been reported in all three basins in the northern plains of Mars: Acidalia Planitia (e.g., Hemmi & Miyamoto 2018; Oehler & Allen 2010) (Fig. 13a), Chryse Planitia (e.g. Brož et al. 2019; Komatsu et al. 2016; Martínez-Alonso et al. 2011) and Utopia Planitia (e.g., McGowan 2011; Skinner & Tanaka 2007) (Fig. 13b). These basins are at a lower elevation than the southern highlands (Fig. 2d) and form a depositional trap for the sediments eroded from the highlands by the valley networks and outflow channels (e.g., Tanaka et al. 2005). At other sites, such as Candor and Coprates Chasmata (Okubo 2016) and Terra Sirenum (Hemmi & Miyamoto 2017) the pitted cones are located in sedimentary fill in other smaller basins. Sedimentary volcanism is thought to mobilise the sediments accumulated in these various basins to form cones and flows. In Acidalia the mud-volcanoes are also associated with polygonal patterns tens of kilometres in diameter, which could represent fluid expulsion pathways (Allen et al. 2013). Fractures and spatially associated pitted cones are also interpreted to be a result of sedimentary volcanism in the Galaxias region (Fig. 13c) (Gallagher et al. 2018) and sedimentary volcanism is a common interpretation of such enigmatic assemblages of km-scale polygonal fractures, ridges and aligned-cones (e.g., De Toffoli et al. 2018; Ivanov et al. 2014; Salvatore & Christensen 2014; Skinner & Mazzini 2009).
<figure><img src="image_30.png" /></figure>

Figure 13: Examples of evidence for sedimentary volcanism on Mars. (a) Pitted cones in Acidalia Planitia with a fracture-junction visible top-left. CaSSIS image MY35_009369_141_0 with channels NIR, PAN and BLU mapped to RGB. (b) Pitted cones and ridges in southern Utopia Planitia. HiRISE image ESP_053364_2135. (c) Aligned cones in Galaxias region, CTX image B05_011746_2156. Image credits: NASA/JPL/UofA/MSSS/ESA/Unibe.

Present-day detections of methane in the martian atmosphere are still heavily debated (e.g., Webster et al. 2018 and references therein), but the discussion of methane has reignited interest in finding and analysing the surface expression of sedimentary volcanism on Mars, as methane gas can be a key driver of sedimentary volcanism on Earth. However, many of these features are estimated as being Amazonian in age or older (Ivanov et al. 2014; Komatsu et al. 2016; Oehler & Allen 2010; Senthil Kumar et al. 2019), and in general are older than gullies and the most recent periglacial and glacial landforms, which means they cannot be *directly* linked to present-day atmospheric methane detections. Sedimentary volcanism implies a source of pressurised liquid water underground whose supply could be engendered by a wide-range of processes including: the impingement of the cryosphere on the hydrosphere due to long-term climate-shifts, destabilisation of clathrates (e.g., De Toffoli et al. 2019) via localised increases in geothermal heat flux or seismic/tectonic/magmatic activity, or simply diapirism of trapped saturated sediments (see summary in Skinner & Tanaka 2007).

Sedimentary volcanism and by association liquid water, is not the only hypothesis for the generation of this suite landforms. Mud-volcanoes could easily be confused with magmatic cones, especially as the behaviour of mud is not well-constrained under martian surface conditions (Brož et al. 2020) and similarly the associations of cones, ridges and fractures have been attributed to a range of processes, e.g. volcanic (e.g., Fagents et al. 2002; Frey & Jarosewich 1982; Ghent et al. 2012), tsunamigenic (Costard et al. 2017) and glacial (e.g., Guidat et al. 2015; Kargel et al. 1995; Rossbacher & Judson 1981). These alternates also imply liquid water, as the volcanic hypothesis requires volatile enrichment to cause the explosive eruptions making the cones (lava-ice or lava-water interaction). The alternates generally imply short-lived episodes of water generation, as opposed to the long-lived aquifer implied in sedimentary volcanism. Liquid water is consistent with the morphology of these features and it is difficult to assess if the sub-surface conditions of Mars could have been conducive to their formation via sedimentary volcanism, as we know so little about the internal structure of Mars (cf. Section 3.2).

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1039

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1057
1058 **3.12 Liquid water related morphologies associated with impacts**
1059 Hypervelocity impacts by bolides into planetary surfaces deposit huge amounts of energy into the target
1060 body, of which most is converted into heat (Melosh 1989). The lack of substantial atmosphere on Mars
1061 means that bolides reach the surface with greater energy than they do on Earth. As discussed in
1062 Section 2 we know that there are substantial amounts of water ice in the near-surface of the martian
1063 crust and potentially a deep cryosphere, and so the heat deposited into the crust by impacts could
1064 generate liquid water by melting. This could be melting immediately during the impact in the cavity and
1065 the ejecta (e.g., Boyce et al. 2012; Newsom 1980; Weiss & Head 2016), vaporisation of the ice could
1066 inject vapour into the surrounding atmosphere and create a transient and localised hydrological cycle
1067 (e.g., Kite et al. 2011; Segura 2002; Segura et al. 2008; Steakley et al. 2019) and finally the thermal
1068 anomaly could maintain hydrothermal circulation within the crust for thousand or millions of years post
1069 impact (e.g., Abramov & Kring 2005; Barnhart et al. 2010; Osinski et al. 2013; Rathbun & Squyres 2002).
<figure><img src="image_31.png" /></figure>

<figure><img src="image_32.png" /></figure>

<figure><img src="image_33.png" /></figure>

<figure><img src="image_34.png" /></figure>

<figure><img src="image_35.png" /></figure>

<figure><img src="image_36.png" /></figure>
1071 Figure 14: Crater-related liquid water morphologies. (a) Pits on the floor of <10 Ma Tooting Crater (Mouginis-Mark & Boyce 2012). CTX image P01_001538_2035. (b) Double-layered ejecta crater in northern plains, with lobate margins and radiating grooves (Weiss & Head 2014). Top HRSC image h2878_0000 and bottom CTX image P20_008833_2149. (c) Multilayer ejecta crater in southern hemisphere with channels incised into its ejecta blanket (Mangold 2012). Top, HRSC image ha313_0000, bottom CTX image B16_015984_2149. (d) Hale Crater with channels incising the ejecta. Top HRSC images h2526_0001 and h0533_0000, bottom CTX image P03_002220_1418. (e) Lyot Crater with clastic polygonal networks in the ejecta. Top THEMIS day IR controlled mosaic image from the USGS and bottom HiRISE image ESP_016985_2315. (f) Mojave Crater with alluvial fans on the inner craters slopes. Top HRSC images hd605_0000 and h1991_0000 and bottom HiRISE image ESP_012834_1875. Image credits: NASA/JPL/MSSS/ESA/DLR/UofA.

1082 Geomorphological evidence for these processes in the Amazonian epoch is abundant. Craters only a few
1083 million years old display pits on their floors (Fig. 14a) and within their ejecta materials interpreted to be a
1084 result of volatile-release as steam (Boyce et al. 2012; Morris et al. 2010; Mouginis-Mark & Boyce 2012;
1085 Tornabene et al. 2012). These pits can be associated with incised channels, alluvial fans (Fig. 14f) and
1086 debris lobes emanating from high points inside the crater cavity, thought to represent the flow of water-
1087 rich materials immediately post-impact, e.g. alluvial fans in Tooting and Mojave Craters both dated to <
1088 10 Ma (Goddard et al. 2014; Morris et al. 2010; Williams & Malin 2008). Channels incising into and
1089 emanating from ejecta blankets are used as evidence for water escaping from the melting cryosphere
1090 beneath (Fig. 14c,d) (El-Maarry et al. 2013; Harrison et al. 2010; Jones et al. 2011; Mangold 2012; Weiss
1091 et al. 2017). Clastic polygonal networks in the ejecta of Lyot Crater (Fig. 14e) even suggest that water in
1092 the ejecta blanket could result in longer-term ice-segregation processes (Brooker et al. 2018) and conical
1093 landforms in the ejecta flows at Hale Crater may also be suggestive of longer-term ice-loss processes
1094 (Conway et al. 2019a).

1095 Lobate ejecta morphologies (e.g., multi-layer ejecta, double-layer ejecta - Fig. 14b) are interpreted to be
1096 the result of fluidisation and are observed almost exclusively at the mid- to high-latitudes (e.g., Barlow
1097 2006; Barlow 2005; Barlow & Perez 2003; Costard 1989; Mouginis-Mark 1981; Mouginis-Mark 1979).
1098 Numerical models generally find this morphology to be consistent with either melting or a volatile-rich
1099 layer (e.g., Baratoux et al. 2002; Oberbeck 2009; Weiss & Head 2014; Weiss & Head 2013). However,
1100 other modelling studies find that the fluidisation is unrelated to water and could be due to increased
1101 fragmentation (Rager et al. 2014) or the presence of an erodible surface (Wada & Barnouin-Jha 2006) –
1102 yet both require preconditioning of the surface by water or water ice to explain the latitudinal
1103 distribution.

1104 Spectral evidence for hydrothermal alteration is common in craters prior the Amazonian, but is lacking in
1105 Amazonian craters (Turner et al. 2016). Similarly, evidence for crater-bound paelolakes, which could
1106 result from hydrothermal circulation are lacking particularly from the late Amazonian epoch (e.g., Cabrol
1107 1999).

1108 Craters also host numerous fluvial channels or fans that cannot necessarily directly be linked to the
1109 impact event itself but could be formed as a result of the impact's topography (microclimate, and/or
1110 sink) or its interaction with groundwater. Examples include: regional drainage networks of Lyot Crater
1111 linked to surface ice accumulation and melt in a climatic microenvironment (Dickson et al. 2009), valleys
draining into craters' central pit (Peel & Fassett 2013), and alluvial fans dated to around ~1 Ga resulting
from prolonged runoff (e.g., Grant & Wilson 2011; Kite et al. 2019; Kite et al. 2017; Mangold et al. 2012).

In summary, impact crater related morphologies in the late Amazonian represent very transient melting
of ground ice mobilised by the impact event, with more substantial evidence for sustained liquid water
environments tending to come from the early Amazonian. In essence, most of these landform
interpretations are consistent with our expectations for the existence of liquid water on Mars, as impacts
are almost intuitively expected to cause melting if they interact with ice, but some of the details of
exactly how much water and its origin are more the focus of debate.

### 3.13 Other evidence for liquid water in the Amazonian

Other morphological evidence for liquid water dates to the early Amazonian with both valley networks
and outflow channels having been reported to have been active at this time. It is generally acknowledged
that water in such systems was short-lived and ephemeral in contrast to the Noachian valley networks
(e.g., Hauber et al. 2013; Wilson et al. 2016). For example, shallow snowmelt has been proposed to
explain valleys in Gorgonum and Newton basins (Howard & Moore 2011), for Kārūn Valles (Adeli et al.
2016) and valley networks on the flanks of martian volcanos (Fassett & Head III 2008). A liquid water
origin for these landforms represents the consensus even if its source is debated and varied.

Some of the martian outflow channels are also dated to the Amazonian (particularly those on the flanks
of the martian volcanos) and their interpretation is less consensual. A detailed discussion of the
competing hypotheses can be found in Chapters LEVERINGTON (pages xxx-xxx) and GALLAGHER (pages
xxx-xxx). In brief the crux of the debate is whether these outflow morphologies can be ascribed to
outpourings of water or mud triggered by lava-ice interaction (e.g., Balme et al. 2011; Burr 2002; Vijayan
& Sinha 2017; Wilson & Head 2002) or deep aquifer sources (e.g., Hanna & Phillips 2006; Marra et al.
2015; Marra et al. 2014), or low viscosity lava (Leverington 2011; Leverington 2009). The debate arises
from the observation that these outflow channels are covered in basalt and not apparently sedimentary
facies, as might be expected from outpourings of liquid water.

## 4 Synthesis and outlook

The reason that the debate around the evidence for liquid water in Mars is still a hot topic is because the
search is motivated by not only scientific curiosity, but can be framed in the larger question of whether
life exists or arose elsewhere in the Solar System. In the search for present or past life beyond our
planet, finding evidence of present or past liquid water on Mars has been a preoccupation of planetary
scientists since the beginning of orbital observations and has been the reason that various space
agencies have sent missions to Mars (e.g., ESA's ExoMars missions; Kereszturi et al. 2016; Vago et al.
2017; and NASA's motto "follow the water"; Hubbard et al. 2002). Without a protective atmosphere the
surface of Mars is bathed in ionising radiation and solar particles, making it a particularly challenging
surface environment for life at the present-day (e.g., Teodoro et al. 2018). However, finding evidence
that liquid water has appeared periodically at the surface could indicate that Mars has subsurface water
reservoirs that could be harbours for extant life. Recent activity of liquid water also bolsters the
probability of liquid water activity in the past and the longevity of habitable conditions. The presence of
habitable conditions also increases the chance that the space probes we send to Mars could contaminate
these likely uninhabited habitats (Cockell et al. 2012) and out of this concern have arisen planetary
protection regulations (National Aeronautics and Space Administration 2005). These regulations
1153 determine the level of sterilisation to which a spacecraft must be submitted to reduce the likelihood of
1154 cross-contamination. This effectively prohibits the direct exploration of zones that could have liquid
1155 water, for example sites with RSL and gullies are subject to these rules (Kminek et al. 2010; Rummel et al.
1156 2014). Water is also an important resource for future human investigations, even in its solid state (e.g.,
1157 Heldmann et al. 2014) and hence finding zones where relatively pure water ice would be easily
1158 accessible has attracted some recent attention (e.g. SWIM project; Perry et al. 2019).

1159 While none of the features described here have been proven beyond all doubt to require the activity of
1160 liquid water, in the aggregate it is challenging to explain all of them without any liquid water on an ice-
1161 rich planet that has straddled the triple point of water for all of the Amazonian. Liquid water is more
1162 generally agreed to be responsible for early Amazonian features including valleys, channels, alluvial fans
1163 and deltas. Most of which are thought to originate from punctual or rare events. However, a consensus
1164 view has not been reached over the role, or not, of sedimentary volcanism in building landforms over
1165 this period. In the late Amazonian to present-day, the role of water is much less certain, with practically
1166 all landforms interpreted as being linked to liquid water having substantial alternate hypotheses. The
1167 exceptions are solifluction lobes and pingos where to-date there is no alternate hypothesis to those
1168 involving liquid water has been proposed. The evidence for sub-glacial melting in the Late Amazonian is a
1169 topic that has only recently been opened and will likely receive more attention as more observations are
1170 collected. For presently active features, the focus is currently on processes that do not involve liquid
1171 water, but the behaviour of certain downslope flows means that liquid water has not been entirely ruled
1172 out. RSL continue to be the most likely of presently active surface features to be linked to liquid water.
1173 For the reasons outlined above it is important to continue studying these landforms and continually
1174 assess the balance of evidence as new data are collected.

1175 The long-lived and ongoing debate over the presence of recent liquid water at the martian surface raises
1176 a couple of important questions for planetary science:

1177 - What level of knowledge do we need to determine (with a reasonable level of confidence)
1178 whether liquid water could be present at the surface?

1179 - Are we being misled by Earth analogues and our *a priori* knowledge of terrestrial
1180 geomorphology?

1181 For the existence of present-day liquid water, RSL are the primary focus of debate, although periodically
1182 other features are proposed (Bhardwaj et al. 2017; Heyer et al. 2019; Kossacki & Markiewicz 2010;
1183 Kreslavsky & Head 2009). Progress on RSL is being hindered by:

1184 - Lack of high temporal resolution monitoring (imaging, SAR, IR and VIS-NIR spectral) at the spatial
1185 resolution of CaSSIS-HiRISE, observations at different times of day or even better continuous
1186 monitoring

1187 - Lack of knowledge of the near-surface atmosphere apart from the handful of places where we
1188 have landed missions with meteorology instruments. This limits our knowledge on the influence
1189 of micro-environments in the steep topography where RSL are found.

1190 - Lack of knowledge of the basic attributes of the martian subsurface, including the crustal heat
1191 flow, porosity, permeability of the regolith with depth, the presence and structure of a potential
1192 cryo-hydro-sphere.

1193 Debates over water in the recent past also depend on the above lacunae and we can also add:
1194 - Limits/uncertainties in our ability to model the martian climate at the present-day are magnified
1195 when applied to the past and limit our ability to understand past environments. Hence, we are
1196 reliant on geomorphic evidence for informing us about the activity of liquid water at the surface,
1197 which in and of itself is reliant on our knowledge of what processes form landscapes on Earth.

1198 The existing hardware in orbit around Mars and on its surface combined with their operational
1199 constraints and observational modes cannot resolve these problems. New hardware could address some
1200 of these issues, for example in orbit: a CaSSIS-HiRISE imager on a spacecraft with a lower orbital
1201 inclination to provide higher frequency imaging and time of day coverage (McEwen et al. 2012), high-
1202 resolution (10 m/pix) thermal imager to better constrain thermal inertia and possibly the temperature of
1203 shallow brine, Synthetic Aperture Radar (SAR) with high resolution (18 m/pix) orbital radar to detect
1204 shallow ice or liquid water, or an imaging spectrometer capable of distinguishing ices and their structure
1205 at high resolution. Surface measurements could include: a network of heat flow probes to determine
1206 how heat flow varies over Mars and better constrain its value, to measure the volume of subsurface ice
1207 (Grimm & Stillman 2015), ground penetrating radar (GPR) to detect the structure of segregated ice and
1208 possibly any shallow aquifers, a transient electromagnetic sounder (TEM) capable of detecting aquifers
1209 to a depth of a few kilometres (e.g., Stamenkovic et al. 2020), or use of the magnetotelluric (MT) method
1210 to passively sense deep conductive bodies and is capable of detecting a subsurface aquifer within the
1211 first 5 km of the crust (e.g., Grimm et al. 2020; Grimm 2002).

1212 Our knowledge from terrestrial geomorphology inevitably leaves us with a biased view of the landscapes
1213 and landforms we find on Mars. Not only that, but even on Earth multiple processes can lead to similar
1214 looking landforms "equifinatlity" - a problem usually resolved on Earth by employing a multidisciplinary
1215 investigation. For example, hummocky landforms can be the result of rock avalanches or be the vestiges
1216 of glacial terminal moraines (e.g., McColl et al. 2019; McColl & Davies 2011). These conundra are
1217 resolved using sedimentology, stratigraphy, geophysical techniques and cosmogenic or other dating
1218 techniques, in combination with field and remote sensing observations (e.g., images, digital terrain
1219 models). On Mars we only have remote sensing data to rely on and sometimes even with images,
1220 spectral data, radar and digital terrain models the interpretation can remain ambiguous. We are also
1221 confronted with an environment where the processes acting on it could, and should, be very different to
1222 those on Earth. For example, the lack of plate tectonics on Mars means that evidence of surface
1223 processes is preserved longer at the surface than we are used to on Earth, meaning landforms with a
1224 "fresh" appearance could date to hundreds of thousands if not millions of years. The generation of relief
1225 on Earth is dominated by the continuous action of plate tectonics and its associated volcanism, whereas
1226 on Mars the relief is created by stochastic impact events and in the past the development of outsized
1227 volcanoes and fault-systems. This means the overall structure of the landscapes of the two planets is
1228 fundamentally different starting at the continental scale.

1229 Despite these differences, as demonstrated in Section 3, remarkably similar landscapes and landforms
1230 are found on the two planets at scales of kilometres to metres. Terrestrial analogues have therefore
1231 helped us to deduce what kinds of processes are active to produce these landscapes. A philosophical
1232 analysis of the utility of Earth analogues for informing planetary investigations is beyond the scope of
1233 this work, but interested readers are referred to the analyses of Baker (2014). Nevertheless, whether or
1234 not "something else" can produce the same landscapes and landforms as liquid water is unknown and is
1235 one of the core uncertainties stemming from our use of terrestrial analogues. For example, for martian
1236 gullies we know the flows at the present-day are fluidised to the same degree as water-saturated debris
1237 flows on Earth (de Haas et al. 2019), but we do not know if sublimating CO₂ could mimic this fluidisation despite being controlled by different physics (de Haas et al. 2019).

1239 Another method of addressing the range of environmental conditions under which water (or brine) can result in landform development is to perform controlled laboratory experiments. Such experiments need to have two complementary outlooks. Firstly, experiments to address very basic physical questions (how fast can brines form by deliquesce in presence of Mars sediments?) in order to test our parameterisation and understanding of the physical laws underpinning the kinetics of individual processes. The aim is to understand the basis physics in a realistic transient environmental setting for Mars. Additionally, understand the metastability of brine can be very important as brine can remain liquid for an extended period of time (e.g., Gough et al. 2011; Toner et al. 2014b; Primm et al. 2017; Primm et al. 2019).
1240
1241
1242
1243
1244
1245
1246
1247 Secondly, more complex simulations, analogous to flume experiments, where processes can be transient and interlinked in order to address the open questions in an empirical fashion. For example, laboratory work can be used to compare the action of water as a fluidising agent to that of gas supported flows.
1248
1249
1250 Particularly in this second case, care needs to be taken to appropriately scale the experiments and consider the limits of their application to Mars (Paola et al. 2009).
1251

1252 5 Conclusions

1253 If we used geomorphological arguments in isolation the case for liquid water on Mars would be hard to 1254 counter. The main argument against liquid water is that martian environmental conditions should not 1255 allow its production, so either we have some fundamental misunderstanding of the martian 1256 surface/subsurface conditions at the present and/or in the recent past, or some other agent(s) is(are) 1257 acting like liquid water does on Earth. For some of the landforms examined in this chapter alternative 1258 processes have been identified, but for others, namely "solifluction" lobes and pingo-like mounds, liquid 1259 water remains the only suggested agent. The lack of agent does not imply one does not exist, but it could 1260 remain to be elucidated. We maintain that liquid water is a viable contender to explain the present-day 1261 activity of RSL, the recent formation of gullies, periglacial landforms and subglacial landforms and the 1262 formation of certain impact-related features and sedimentary volcanism in the deeper past. The 1263 implication being that, although a lot less active than in the past, liquid water remains an active 1264 geomorphological agent on Amazonian Mars. In order to consolidate this position, we argue that the 1265 most progress can be made in the short-term by focussing research efforts on experimental work 1266 informed by remote sensing observations and in the long-term by better equipped Mars-missions in the 1267 future. Additionally, sending a small mission to determine, if Mars still possesses deep groundwater 1268 would be greatly beneficial as it can serve as an end member for the likelihood of water to have survived 1269 in the ground for 100 Ma to Ga.

1270 6 References cited

1271 Abotalib, A.Z. & Heggy, E., 2019. A deep groundwater origin for recurring slope lineae on Mars. *Nature Geoscience*, 12(4), p.235–241.

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