TThe Climate of Early Mars
Robin D. Wordsworth , Harvard Paulson School of Engineering and Applied Sciences, HarvardUniversity, Cambridge, MA 02140, USA Department of Earth and Planetary Sciences, Harvard University, Cambridge,MA 02140, USAAnnual Review of Earth & PlanetarySciences 2016. 44:1–31This article’s doi:10.1146/xxxxxxxCopyright © Keywords mars, paleoclimate, atmospheric evolution, faint young Sun,astrobiology
Abstract
The nature of the early Martian climate is one of the major unansweredquestions of planetary science. Key challenges remain, but a new waveof orbital and in situ observations and improvements in climate mod-eling have led to significant advances over the last decade. Multiplelines of geologic evidence now point to an episodically warm surfaceduring the late Noachian and early Hesperian periods 3-4 Ga. The lowsolar flux received by Mars in its first billion years and inefficiency ofplausible greenhouse gases such as CO means that the steady-stateearly Martian climate was likely cold. A denser CO atmosphere wouldhave caused adiabatic cooling of the surface and hence migration ofwater ice to the higher altitude equatorial and southern regions of theplanet. Transient warming caused melting of snow and ice deposits anda temporarily active hydrological cycle, leading to erosion of the val-ley networks and other fluvial features. Precise details of the warmingmechanisms remain unclear, but impacts, volcanism and orbital forcingall likely played an important role. The lack of evidence for glaciationacross much of Mars’ ancient terrain suggests the late Noachian surfacewater inventory was not sufficient to sustain a northern ocean. Whilemainly inhospitable on the surface, early Mars may nonetheless havepresented significant opportunities for the development of microbial life. a r X i v : . [ a s t r o - ph . E P ] J un . INTRODUCTION With the exception of Earth, Mars is the Solar System’s best-studied planet. Since the first
Mariner flybys in the 1960s, Mars has been successfully observed by a total of 11 orbitersand 7 landers, four of them rovers. Combined with ongoing observations from Earth, thishas allowed a uniquely comprehensive description of the Martian atmosphere and surface.However, despite the wealth of data obtained, fundamental mysteries about Mars’ evolutionremain. The biggest mystery of all is the nature of the early climate: 3-4 Ga Mars shouldhave been freezing cold, but there is nonetheless abundant evidence that liquid water flowedacross its surface.Unlike Earth, Mars lacks plate tectonics, global oceans and a biosphere . One of thehappy outcomes of this is that its ancient crust is incredibly well-preserved, allowing awindow to epochs as early as 3-4 Ga across large regions of the surface (Nimmo & Tanaka2005). This antiquity is hard to imagine from a terrestrial perspective. By way of compar-ison, we can imagine the advantages to Precambrian geology if most of Asia consisted oflightly altered terrain from the early Archean, when life was first emerging on Earth. Marsprovides us a glimpse of conditions during the earliest stages of the Solar System on a bodythat had an atmosphere, at least episodic surface liquid water, and in some locales, surfacechemistry conducive to the survival of microbial life (e.g., Grotzinger et al. 2014).There are many motivations for studying the early climate of Mars. The first is simplythat it is a fundamentally interesting unsolved problem in planetary science. Another majormotivation is astrobiological — if we can understand how the Martian climate evolved, wewill have a better understanding of whether life could have ever flourished, and where tolook for it if it did. Studying Mars also has the potential to inform us about the evolutionof our own planet, because many of the processes thought to be significant to climate onearly Mars (e.g. volcanism, impacts) have also been of major importance on Earth. Finally,in this era of exoplanet science, Mars also represents a test case that can inform us aboutthe climates of small rocky planets in general.The aim of this review is to provide a general introduction to the latest research on theearly Martian climate. We begin by discussing highlights of the geologic evidence for analtered climate on early Mars, focusing on the extent to which the observations are consistentwith episodic warming vs. a steady-state warm and wet climate. Next, we discuss theexternal boundary conditions on the early climate (namely the solar flux and Martian orbitalparameters) and the constraints on the early atmospheric pressure. We also review previousone-dimensional radiative-convective modeling of the effects of key processes (atmosphericcomposition, meteorite impacts and volcanism) on surface temperature. Finally, we discussrecent three-dimensional climate modeling of early Mars by a number of groups that hasincreased our understanding of cloud and aerosol processes and the nature of the earlywater cycle. It is argued that future progress will require an integrated approach, wherethree-dimensional climate models are compared with the geologic evidence on both globaland regional scales. While the possibility of life on Mars today still cannot be entirely ruled out, the absence of abiosphere sufficient to modify the surface substantially is clear. . GEOLOGIC EVIDENCE FOR LIQUID WATER ON EARLY MARS
With a few important exceptions, all our current information on the early martian climatecomes from surface geology. Martian geologic data is derived from a combination of passiveand active orbital remote sensing and in situ analysis. The oldest and best studied aspectof Martian geology is the surface geomorphology. In recent years, the geomorphic data havebeen supplemented by global maps of surface mineralogy derived from orbiters and detailed in situ studies of several specific regions by the NASA rover missions.Figure 1 (top left) summarizes the basic features of the Martian surface. The four mostimportant large-scale features are the north-south dichotomy, the Tharsis bulge, and theHellas and Argyre impact craters. Because Martian topography plays a major role in theplanet’s climate and hydrological cycle, understanding when these features formed is vital.The relative ages of surface features and regions on Mars can be assessed via analysis oflocal crater size-frequency distributions (crater statistics) (Tanaka 1986; Tanaka et al. 2014).Absent geochronology data, absolute dating of Martian surface units relies on impactor fluxmodels and hence is subject to considerable uncertainty. (It is standard to categorize Martian terrain into three time periods (Fig. 2): the mostmodern Amazonian ( ∼ − . ∼ . − . ∼ . − . The valley networks are the single most important piece of evidencein favor of a radically different climate on early Mars. Like many drainage basins onEarth and in contrast with the later Hesperian-period outflow channels, Martian valleynetworks are dendritic (branching) with tributaries that begin near the peaks of topographicdivides. This geomorphology strongly suggests an origin due to a hydrological cycle drivenby precipitation (as rain or snow) (Craddock & Howard 2002; Mangold et al. 2004; Stepinski& Stepinski 2005; Barnhart, Howard & Moore 2009; Hynek, Beach & Hoke 2010; Matsubara,Howard & Gochenour 2013) rather than e.g., groundwater sapping (Squyres & Kasting 1994)or basal melting of thick icesheets (Carr & Head 2003). ••
With a few important exceptions, all our current information on the early martian climatecomes from surface geology. Martian geologic data is derived from a combination of passiveand active orbital remote sensing and in situ analysis. The oldest and best studied aspectof Martian geology is the surface geomorphology. In recent years, the geomorphic data havebeen supplemented by global maps of surface mineralogy derived from orbiters and detailed in situ studies of several specific regions by the NASA rover missions.Figure 1 (top left) summarizes the basic features of the Martian surface. The four mostimportant large-scale features are the north-south dichotomy, the Tharsis bulge, and theHellas and Argyre impact craters. Because Martian topography plays a major role in theplanet’s climate and hydrological cycle, understanding when these features formed is vital.The relative ages of surface features and regions on Mars can be assessed via analysis oflocal crater size-frequency distributions (crater statistics) (Tanaka 1986; Tanaka et al. 2014).Absent geochronology data, absolute dating of Martian surface units relies on impactor fluxmodels and hence is subject to considerable uncertainty. (It is standard to categorize Martian terrain into three time periods (Fig. 2): the mostmodern Amazonian ( ∼ − . ∼ . − . ∼ . − . The valley networks are the single most important piece of evidencein favor of a radically different climate on early Mars. Like many drainage basins onEarth and in contrast with the later Hesperian-period outflow channels, Martian valleynetworks are dendritic (branching) with tributaries that begin near the peaks of topographicdivides. This geomorphology strongly suggests an origin due to a hydrological cycle drivenby precipitation (as rain or snow) (Craddock & Howard 2002; Mangold et al. 2004; Stepinski& Stepinski 2005; Barnhart, Howard & Moore 2009; Hynek, Beach & Hoke 2010; Matsubara,Howard & Gochenour 2013) rather than e.g., groundwater sapping (Squyres & Kasting 1994)or basal melting of thick icesheets (Carr & Head 2003). •• The Climate of Early Mars 3 alley networks are rare on Hesperian and Amazonian terrain but common on Noachianterrain, where they are predominantly seen at equatorial latitudes between 60 ◦ S and 10 ◦ N(Milton 1973; Carr 1996; Hynek, Beach & Hoke 2010). The largest networks are huge,extending thousands of kilometers over the surface in some cases (Howard, Moore & Irwin2005; Hoke, Hynek & Tucker 2011). Landform evolution models suggest minimum formationtimescales for the valley networks of 10 to 10 years under climate conditions appropriateto arid regions on Earth (Barnhart, Howard & Moore 2009; Hoke, Hynek & Tucker 2011).( When liquid water carves valley networks on a heavily cratered terrain,ponding and lake formation inside craters is a natural outcome. For sufficiently high flowrates, crater lakes will breach their rims, forming open lakes that are integrated in a largerhydrological network. Both the ratio of watershed area to lake area (drainage ratio) foreach lake and the ratio of open to closed basin crater lakes in total give important cluesas to the nature of the Noachian water cycle. In general, low drainage ratios and a largenumber of open crater lakes indicate high precipitation rates and a wet climate (Fassett &Head 2008; Barnhart, Howard & Moore 2009).As might be expected, analysis of the Noachian southern highlands has revealed abun-dant evidence for crater lakes interlinked with the valley networks (Cabrol & Grin 1999;Fassett & Head 2008). However, closed-basin lakes greatly outnumber open-basin lakes.This suggests that a very wet climate or periodic catastrophic deluges due to e.g. impact-driven steam greenhouses (see Section 3.4) were not responsible for their formation (Irwinet al. 2005; Barnhart, Howard & Moore 2009). Open-basin lake drainage ratios stronglyvary with location, with wetter formation conditions indicated in Arabia Terra and northof Hellas (Terra Sabaea) (Fassett & Head 2008).Striking evidence of in situ fluvial erosion was found by NASA’s Curiosity rover inthe form of conglomerate outcrops at Gale Crater (Williams et al. 2013). Morphologically,the conglomerates are remarkably similar to sediment deposits found on Earth (see Fig. 1).However, chemical analysis of the outcrops suggested low chemical alteration of the materialby water (Williams et al. 2013). Indeed, global analysis of aqueous alteration productson the Martian surface suggests a predominance of juvenile or weakly modified minerals(Tosca & Knoll 2009). In addition, most of the minerals in open-basin crater lakes observedfrom orbit lack evidence of strong in situ chemical alteration (Goudge et al. 2012). Thissuggests that the flows responsible for eroding the late Noachian and Hesperian surfacewere relatively short-lived.
The most evocative (and controversial) claim to have come out ofgeomorphic studies of the ancient surface is that Mars once possessed a northern oceanof liquid water. The original argument proposed to support this is that various geologiccontacts in the northern plains resemble ancient shorelines (Parker et al. 1993; Head et al.1999). Several of the putative shorelines show vertical variations of several kilometers, whichis inconsistent with a fluid in hydrostatic equilibrium, although it has been argued thattrue polar wander could have caused surface deformation sufficient to explain this (Perronet al. 2007). More critically, much of the shoreline evidence was found to be ambiguous insubsequent high-resolution imaging studies (Malin & Edgett 1999).More recently, it has been argued that many delta-like deposits, which are assumed tobe of fluvial origin, follow an isostatic line at -2.54 km from the datum across the surface (diAchille & Hynek 2010). If this line does represent a Noachian ocean shoreline, the implica- ions are a) that the earlier proposed shorelines are incorrect and Martian topography wasnot modified by a true polar wander event and b) Mars once had a global equivalent layer(GEL) of around 550 m of surface water. The extent to which a warm and wet scenariofor early Mars with a northern ocean fits the climate constraints and the other geologicevidence is a major focus of the rest of this article. (
The evidence for an at least episodically warmer early Martian climateis not limited to fluvial landforms. At the south pole, the Dorsa Argentea Formation(DAF), a geologic unit dated to the mid-Hesperian, contains a range of features suggestive ofglaciation, including sinuous ridges interpreted as eskers (Fig. 1) and pitted regions that mayhave been caused by basal melting of a thick ice sheet (Howard 1981; Head & Pratt 2001).Around the Argyre and Hellas basins, further glacial landforms such as eskers, lobate debrisaprons and possible moraines and cirques are observed (Kargel & Strom 1992). Dynamicicesheet modeling (Fastook et al. 2012) suggests that polar surface temperature increasesof 25-50 K from Amazonian (modern) values are required before wet-based glaciation ofthe DAF could occur, again suggesting episodically warmer climate conditions on earlyMars. Interestingly, however there is comparatively little evidence for glacial alterationof the surface on Noachian or Hesperian terrain at more equatorial latitudes. This is animportant issue that we return to in Section 4. ((
The morphological evidence for liquid water on early Mars, which has been observed inincreasing detail from the 1960s onwards, has been complemented over the last 10-15 yearsby a new array of geochemical observations from orbit and rover missions. Iron and magne-sium rich phyllosilicates (clays) are found extensively over Noachian terrain (Poulet et al.2005; Bibring et al. 2006; Mustard et al. 2008; Murchie et al. 2009; Carter et al. 2010). Toform, these minerals require the presence of liquid water and near-neutral-pH conditions.Other aqueous minerals such as sulphates, chlorides and silicas are found in more localizedregions of the Noachian and Hesperian crust (Gendrin et al. 2005; Osterloo et al. 2010;Carter et al. 2013; Ehlmann & Edwards 2014).If the phyllosilicates mainly formed on the surface, this would represent evidence infavor of a warm and wet early Martian climate (Poulet et al. 2005; Bibring et al. 2006).Recently, however, it has been argued that in most cases their mineralogy may best representsubsurface formation in geothermally heated, water-poor systems (Ehlmann et al. 2011).Many of the sedimentary phyllosilicates observed on the Martian surface today may nothave formed in situ , but were instead transported to their current locations by later erosionof the crust.At several sites, Al-rich clays such as kaolinite are observed alongside or overlying Fe/Mgclays (Poulet et al. 2005; Wray et al. 2008; Ehlmann et al. 2009; Carter et al. 2015). OnEarth, Al-rich clays overlying Fe/Mg clay in a stratigraphic section is a common featureof wet environments, because iron and magnesium cations from the original minerals arepreferentially leached (flushed) downwards from the topmost layer by water from rain orsnowmelt. This is one possible explanation for the presence of Al-rich clays on Mars (Carteret al. 2015). Another interpretation is more acidic and oxidizing local alteration conditionsin a mainly cold climate, as suggested by the presence of the sulfate mineral jarosite adjacentto Al clays in several regions (Ehlmann & Dundar 2015). ••
The morphological evidence for liquid water on early Mars, which has been observed inincreasing detail from the 1960s onwards, has been complemented over the last 10-15 yearsby a new array of geochemical observations from orbit and rover missions. Iron and magne-sium rich phyllosilicates (clays) are found extensively over Noachian terrain (Poulet et al.2005; Bibring et al. 2006; Mustard et al. 2008; Murchie et al. 2009; Carter et al. 2010). Toform, these minerals require the presence of liquid water and near-neutral-pH conditions.Other aqueous minerals such as sulphates, chlorides and silicas are found in more localizedregions of the Noachian and Hesperian crust (Gendrin et al. 2005; Osterloo et al. 2010;Carter et al. 2013; Ehlmann & Edwards 2014).If the phyllosilicates mainly formed on the surface, this would represent evidence infavor of a warm and wet early Martian climate (Poulet et al. 2005; Bibring et al. 2006).Recently, however, it has been argued that in most cases their mineralogy may best representsubsurface formation in geothermally heated, water-poor systems (Ehlmann et al. 2011).Many of the sedimentary phyllosilicates observed on the Martian surface today may nothave formed in situ , but were instead transported to their current locations by later erosionof the crust.At several sites, Al-rich clays such as kaolinite are observed alongside or overlying Fe/Mgclays (Poulet et al. 2005; Wray et al. 2008; Ehlmann et al. 2009; Carter et al. 2015). OnEarth, Al-rich clays overlying Fe/Mg clay in a stratigraphic section is a common featureof wet environments, because iron and magnesium cations from the original minerals arepreferentially leached (flushed) downwards from the topmost layer by water from rain orsnowmelt. This is one possible explanation for the presence of Al-rich clays on Mars (Carteret al. 2015). Another interpretation is more acidic and oxidizing local alteration conditionsin a mainly cold climate, as suggested by the presence of the sulfate mineral jarosite adjacentto Al clays in several regions (Ehlmann & Dundar 2015). •• The Climate of Early Mars 5 ulfate deposits on Mars are particularly interesting because they require a source (mostlikely volcanic) of sulfur and in some cases indicate acidic and/or saline formation conditions.Sulfates appear primarily, but not exclusively, on Hesperian terrain and may be associatedwith the volcanic activity that formed the basaltic ridged plains (Head, Kreslavsky & Pratt2002; Bibring et al. 2006). The link between the sulfates, volcanism and possible changesin the early climate due to sulfur dioxide (SO ) and hydrogen sulfide (H S) is discussed inSection 3.4.2.( One mineral, carbonate, is conspicuous by its low abundance on the Martian surface(Niles et al. 2013; Ehlmann & Edwards 2014). This is important, because surface carbonateformation should be very efficient on a warm and wet planet with a basaltic crust and CO -rich atmosphere (Pollack et al. 1987). Carbonate formation could have been suppressedby acidic surface conditions caused by dissolution of SO in water (Bullock & Moore 2007;Halevy, Zuber & Schrag 2007). However, globally acidic warm and wet conditions are diffi-cult to justify in the presence of a strongly mafic basaltic regolith, which should effectivelybuffer pH, just like the basaltic seafloor on Earth (Niles et al. 2013). Hence the absenceof surface carbonates is a strong indication that early Mars was either only episodicallywarm, or very dry. Interestingly, carbonates have been discovered in outcrops in the NiliFossae region by CRISM (Ehlmann et al. 2008) and in Gusev crater by the Spirit rover(Morris et al. 2010). They are also seen in some regions where deep crustal material hasbeen excavated by impacts (Michalski & Niles 2010). This indicates that regardless of theearly surface conditions, the deep crust may still have been a major sink for atmosphericCO over time (Section 3.1).
3. FAINT YOUNG SUN, COLD YOUNG PLANET?
The geologic record is unanimous: liquid water substantially modified Mars’ surface duringthe late Noachian. The surface processes that could have created this water, however, arefar from obvious. Two basic facts conspire to make warming early Mars a fiendish challenge:the Martian orbit and the faintness of the young Sun.With a semi-major axis of 1.524 AU, Mars receives around 43% of the solar energy thatEarth does. The rapid dissipation of the nebula during terrestrial planet formation and lackof major configurational changes in the Solar System since the late heavy bombardmentmeans that Mars’ orbital semi-major axis cannot have changed significantly since the lateNoachian. Mars’ orbital eccentricity and obliquity evolve chaotically on long timescales,however, and have probably varied over ranges of 0 – 0.125 and 10 ◦ – 60 ◦ , respectively(Laskar et al. 2004). Although this has little effect on the net annual solar flux, it still hasimportant implications for the early climate, because the time-varying insolation patternis a key determinant of peak summertime temperatures and hence the long-term transportand melting of water ice.The early Sun was less luminous than today because hydrogen burning increases themean molar mass of the core, causing it to contract and heat up. The rate of fusionis strongly dependent on temperature, so this in turn increases a main sequence star’sluminosity over time. This fundamental outcome of stellar physics is supported by detailedsolar models and observations of many nearby stars. As a result, the Sun’s luminosity3.8 Ga was approximately 75% of its present-day value (Gough 1981). One possible wayto avoid this outcome is if the Sun shed large amounts of its mass early on [over 2% in thefirst 2 Gyr; Minton & Malhotra (2007)]. While possible, this is unlikely, because such high ass loss is not observed in any nearby G or K class stars (Minton & Malhotra 2007). Theidea that our Sun must be a unique and unusual star solely because Mars once had surfaceliquid water has not gained widespread acceptance.If we accept the standard orbital and solar boundary conditions, the early Mars climateproblem is now simply understood. Let us assume that in the late Noachian, Mars’ receivedsolar flux was 0 . × / . = 441 . − . If the planetary albedo were zero (i.e.every solar photon intersected by Mars was absorbed), the equilibrium temperature wouldthen be T e = [441 . / σ ] / = 210 K (here σ is the Stefan-Boltzmann constant). Thisimplies a minimum greenhouse effect of around 65 K (around double that of present-dayEarth or 9 times that of present-day Mars) to achieve even marginally warm and wet surfaceconditions. For more realistic planetary albedo estimates, the greenhouse effect required istens of degrees greater still. One seemingly obvious way to invoke a more potent greenhouse effect on early Mars is via adenser atmosphere. But how thick could the early atmosphere have been? Estimating thetotal atmospheric pressure in the late Noachian requires consideration of the major sources(volcanic outgassing and impact delivery) and sinks (escape to space and incorporation ofCO into the crust).Carbon dioxide is generally assumed to have been the dominant constituent of theMartian atmosphere in the late Noachian, as it is today. The outgassing of CO intothe Martian atmosphere with time is a function of the rate of volcanic activity and thechemical composition of the mantle (Grott et al. 2011). The rate of volcanism throughthe pre-Noachian and Noachian is not strongly constrained, but the majority of volatileoutgassing in Mars’ history almost certainly occurred in these periods (Grott et al. 2011).Regarding the chemical composition, CO outgassing is strongly dependent in particularon the mantle oxygen fugacity ( f O ). Analysis of Martian meteorites suggests Mars has amore reducing mantle than Earth, with an f O value between the iron-w¨ustite and quartz-fayalite-magnetite buffers (Wadhwa 2001). Given this, Hirschmann & Withers (2008) esti-mate that between 70 mbar and 13 bar of CO could have been outgassed during the initialformation of the Martian crust, with the lower estimate for the most reducing conditions.During later events (such as the formation of the Tharsis bulge), they estimate that 40 mbarto 1.4 bars could have been outgassed. (The escape rate of CO to space until the Hesperian is highly uncertain. Escape rateswere highest just after Mars’ formation, and the isotopic fractionation of noble gases in theatmosphere indicates the majority of the primordial atmosphere was lost very early (Jakosky& Jones 1997). It has also been argued that all of the initially outgassed CO would havebeen rapidly lost to space by XUV-driven escape before the late Noachian (Tian et al. 2010).However, effective loss requires total dissociation of CO into its constituent atoms. Thechemistry of this process has not been extensively studied and may be somewhat model-dependent (Lammer et al. 2013). Meteorite impacts during accretion remove CO butalso deliver it, with the balance dependent on the model used. In contrast, escape processesoccurring from the Hesperian onwards appear unambiguously ineffective: ion escape, plasmainstability, sputtering and non-thermal processes combined could not have removed morethan a few 100 mbar at most since 4 Ga (Chassefi`ere & Leblanc 2004; Lammer et al. 2013).Although further insights into these processes will be supplied by NASA’s ongoing MAVEN ••
The geologic record is unanimous: liquid water substantially modified Mars’ surface duringthe late Noachian. The surface processes that could have created this water, however, arefar from obvious. Two basic facts conspire to make warming early Mars a fiendish challenge:the Martian orbit and the faintness of the young Sun.With a semi-major axis of 1.524 AU, Mars receives around 43% of the solar energy thatEarth does. The rapid dissipation of the nebula during terrestrial planet formation and lackof major configurational changes in the Solar System since the late heavy bombardmentmeans that Mars’ orbital semi-major axis cannot have changed significantly since the lateNoachian. Mars’ orbital eccentricity and obliquity evolve chaotically on long timescales,however, and have probably varied over ranges of 0 – 0.125 and 10 ◦ – 60 ◦ , respectively(Laskar et al. 2004). Although this has little effect on the net annual solar flux, it still hasimportant implications for the early climate, because the time-varying insolation patternis a key determinant of peak summertime temperatures and hence the long-term transportand melting of water ice.The early Sun was less luminous than today because hydrogen burning increases themean molar mass of the core, causing it to contract and heat up. The rate of fusionis strongly dependent on temperature, so this in turn increases a main sequence star’sluminosity over time. This fundamental outcome of stellar physics is supported by detailedsolar models and observations of many nearby stars. As a result, the Sun’s luminosity3.8 Ga was approximately 75% of its present-day value (Gough 1981). One possible wayto avoid this outcome is if the Sun shed large amounts of its mass early on [over 2% in thefirst 2 Gyr; Minton & Malhotra (2007)]. While possible, this is unlikely, because such high ass loss is not observed in any nearby G or K class stars (Minton & Malhotra 2007). Theidea that our Sun must be a unique and unusual star solely because Mars once had surfaceliquid water has not gained widespread acceptance.If we accept the standard orbital and solar boundary conditions, the early Mars climateproblem is now simply understood. Let us assume that in the late Noachian, Mars’ receivedsolar flux was 0 . × / . = 441 . − . If the planetary albedo were zero (i.e.every solar photon intersected by Mars was absorbed), the equilibrium temperature wouldthen be T e = [441 . / σ ] / = 210 K (here σ is the Stefan-Boltzmann constant). Thisimplies a minimum greenhouse effect of around 65 K (around double that of present-dayEarth or 9 times that of present-day Mars) to achieve even marginally warm and wet surfaceconditions. For more realistic planetary albedo estimates, the greenhouse effect required istens of degrees greater still. One seemingly obvious way to invoke a more potent greenhouse effect on early Mars is via adenser atmosphere. But how thick could the early atmosphere have been? Estimating thetotal atmospheric pressure in the late Noachian requires consideration of the major sources(volcanic outgassing and impact delivery) and sinks (escape to space and incorporation ofCO into the crust).Carbon dioxide is generally assumed to have been the dominant constituent of theMartian atmosphere in the late Noachian, as it is today. The outgassing of CO intothe Martian atmosphere with time is a function of the rate of volcanic activity and thechemical composition of the mantle (Grott et al. 2011). The rate of volcanism throughthe pre-Noachian and Noachian is not strongly constrained, but the majority of volatileoutgassing in Mars’ history almost certainly occurred in these periods (Grott et al. 2011).Regarding the chemical composition, CO outgassing is strongly dependent in particularon the mantle oxygen fugacity ( f O ). Analysis of Martian meteorites suggests Mars has amore reducing mantle than Earth, with an f O value between the iron-w¨ustite and quartz-fayalite-magnetite buffers (Wadhwa 2001). Given this, Hirschmann & Withers (2008) esti-mate that between 70 mbar and 13 bar of CO could have been outgassed during the initialformation of the Martian crust, with the lower estimate for the most reducing conditions.During later events (such as the formation of the Tharsis bulge), they estimate that 40 mbarto 1.4 bars could have been outgassed. (The escape rate of CO to space until the Hesperian is highly uncertain. Escape rateswere highest just after Mars’ formation, and the isotopic fractionation of noble gases in theatmosphere indicates the majority of the primordial atmosphere was lost very early (Jakosky& Jones 1997). It has also been argued that all of the initially outgassed CO would havebeen rapidly lost to space by XUV-driven escape before the late Noachian (Tian et al. 2010).However, effective loss requires total dissociation of CO into its constituent atoms. Thechemistry of this process has not been extensively studied and may be somewhat model-dependent (Lammer et al. 2013). Meteorite impacts during accretion remove CO butalso deliver it, with the balance dependent on the model used. In contrast, escape processesoccurring from the Hesperian onwards appear unambiguously ineffective: ion escape, plasmainstability, sputtering and non-thermal processes combined could not have removed morethan a few 100 mbar at most since 4 Ga (Chassefi`ere & Leblanc 2004; Lammer et al. 2013).Although further insights into these processes will be supplied by NASA’s ongoing MAVEN •• The Climate of Early Mars 7 ission, it currently appears that a dense late Noachian atmosphere can only have beenremoved subsequently by surface processes.The most efficient potential sink for atmospheric CO at the surface is carbonate forma-tion. As we have discussed, surface carbonates are rare on Mars. However, the discovery ofcarbonates in exhumed deep crust suggests that the possibility of a large subsurface reser-voir cannot be discounted (Michalski & Niles 2010). Recently, it has been argued based onorbital observations that this reservoir is unlikely to allow more than a 500 mbar CO at-mosphere during the late Noachian (Edwards & Ehlmann 2015). However, sequestration byhydrothermal circulation of CO in deep basaltic crust is a very poorly understood processeven on Earth (e.g., Brady & G´ıslason 1997), so caution is still required when extrapolat-ing the known carbonate reservoirs. It will be hard to constrain the late Noachian carbonbudget definitively until we can send missions to Mars (robotic or human) that are capableof drilling deep into the subsurface.Finally, one independent constraint on atmospheric pressure comes from the size dis-tribution of craters on ancient terrain. In a thick atmosphere smaller impactors burn upbefore they reach the surface, so observation of the smallest craters leads to an upper limiton atmospheric pressure. Recent analysis of the Dorsa Aeolis region near Gale crater usingthis technique has led to an approximate upper limit of 0.9-3.8 bars on atmospheric pressure3.6 Ga (Kite et al. 2014).To summarize, many aspects of Mars’ atmospheric evolution are highly uncertain. Itis likely that Mars had a thicker CO atmosphere in the late Noachian. This atmospherecould have been as dense as 1-2 bar, but likely no more than this. If the early atmospherewas denser than about 0.5 bar, it cannot have all escaped to space and the difference willnow be buried in the deep crust as carbonate. Several recent studies have suggested thatthis reservoir may be small, but the observational search for carbonate deposits on Marsshould continue, along with theoretical study of the interaction between atmospheric CO and pore water in deep Martian hydrothermal systems. Greenhouse
Constraining the early atmospheric CO content is necessary to build a complete picture ofthe Noachian climate, but it is not sufficient. In a seminal paper, Kasting (1991) demon-strated that regardless of the atmospheric pressure, a clear-sky CO -H O atmosphere alonecould not have warmed early Mars. There are two reasons for this. First, CO is an effi-cient Rayleigh scatterer, so in large quantities it significantly raises the planetary albedo.In addition, CO condenses into clouds of dry ice at low temperatures. As surface pressureincreases this leads to a shallower atmospheric lapse rate, reducing the greenhouse effect(Fig. 4). At high enough CO pressures, the atmosphere collapses on the surface completely.This conclusion, which was reached by Kasting using a one-dimensional clear-sky radiative-convective climate model, has recently been confirmed by three-dimensional climate modelsthat include cloud effects (Forget et al. 2013; Wordsworth et al. 2013).(( Although uncertainty remains, the infrared radiative effects of dense CO -dominatedatmospheres are now fairly well understood. CO is opaque across important regions ofthe infrared because of direct vibrational-rotational absorption, particularly due to the ν
667 cm − (15 µ m) bending mode and associated overtone bands. In dense atmospheres,CO , like most gases, also absorbs effectively through collision-induced absorption (CIA).CIA is a collective effect that involves the interaction of electromagnetic radiation with pairs or larger numbers) of molecules. For CO , it occurs due to both induced dipole effectsin the 0-250 cm − region (Gruszka & Borysow 1997), and dimer effects between 1200 and1500 cm − (Baranov, Lafferty & Fraser 2004) (see Fig. 3). Further complications arisefrom the fact that the sub-Lorentzian nature of absorption lines far from their centers mustalso be taken into account for accurate computation of CO absorption in climate models(Perrin & Hartmann 1989).Figure 3 (top) shows the absorption coefficient of CO and other gases at standardtemperature and pressure computed using the CIA and sublorentzian line broadeningparametrization described in Wordsworth, Forget & Eymet (2010), with the two regionsof CIA indicated. In the lower panel, the outgoing longwave radiation (OLR) computed as-suming a 1 bar dry CO atmosphere with surface temperature of 250 K is shown. As can beseen, the CIA causes significant reduction in OLR, but important window regions remain,particularly around 400 and 1000 cm − . The most effective minor greenhouse gases onearly Mars are those whose absorption peaks in these CO window regions. Water vapourabsorbs strongly at low wavenumbers and around its ν band at 1600 cm − , but its molarconcentration is determined by temperature and hence it can only cause a feedback effect onthe radiative forcing of other gases. At low temperature, this feedback is weak. Hence forpure clear-sky CO -H O atmospheres under early Martian conditions, modern radiative-convective models obtain mean surface temperatures of 225 K or less (Wordsworth, Forget& Eymet 2010; Ramirez et al. 2014; Halevy & Head 2014). (
The popular notion that Mars was once warm and wet combined with the impossibility ofwarming early Mars by CO alone has motivated investigation of many alternative warmingmechanisms. For the most part, researchers have used one-dimensional radiative-convectivemodels [either line-by-line or correlated- k ; Goody et al. (1989)] to investigate the earlyclimate. Radiative-convective models allow for temperature variations with altitude onlyand parametrize or neglect the effects of clouds, which strongly limits their accuracy andpredictive power. Nonetheless, their speed and robustness make them invaluable tools forconstraining parameter space and investigating novel effects. ((Over the years, researchers have investigated various greenhouse gas combinations toachieve a warm and wet early Martian climate (Sagan 1977; Postawko & Kuhn 1986; Kast-ing 1997; Haberle 1998; von Paris et al. 2013; Ramirez et al. 2014). Methane might appearto be a promising Martian greenhouse gas given its strong radiative forcing on the present-day Earth. However, it is not an effective warming agent on early Mars because its firstsignificant vibration-rotation band absorbs around 1300 cm − — a region too far from thepeak of the Planck function in the 200-270 K temperature range to cause much change tooutgoing longwave radiation (see Fig. 3). Gases such as ammonia and carbonyl sulphidehave greater radiative potential but lack efficient formation mechanisms and are photochem-ically unstable in the Martian atmosphere, so they cannot have been present long-term inlarge quantities.Other researchers have looked at hydrogen as a greenhouse gas (Sagan 1977; Ramirezet al. 2014). Wordsworth & Pierrehumbert (2013) showed that N -H CIA can cause sig-nificant warming on terrestrial planets even when H is a minor atmospheric constituent,because broadening of the CIA spectrum at moderate temperatures causes absorption toextend into window regions. In a thought-provoking paper, Ramirez et al. (2014) proposed ••
The popular notion that Mars was once warm and wet combined with the impossibility ofwarming early Mars by CO alone has motivated investigation of many alternative warmingmechanisms. For the most part, researchers have used one-dimensional radiative-convectivemodels [either line-by-line or correlated- k ; Goody et al. (1989)] to investigate the earlyclimate. Radiative-convective models allow for temperature variations with altitude onlyand parametrize or neglect the effects of clouds, which strongly limits their accuracy andpredictive power. Nonetheless, their speed and robustness make them invaluable tools forconstraining parameter space and investigating novel effects. ((Over the years, researchers have investigated various greenhouse gas combinations toachieve a warm and wet early Martian climate (Sagan 1977; Postawko & Kuhn 1986; Kast-ing 1997; Haberle 1998; von Paris et al. 2013; Ramirez et al. 2014). Methane might appearto be a promising Martian greenhouse gas given its strong radiative forcing on the present-day Earth. However, it is not an effective warming agent on early Mars because its firstsignificant vibration-rotation band absorbs around 1300 cm − — a region too far from thepeak of the Planck function in the 200-270 K temperature range to cause much change tooutgoing longwave radiation (see Fig. 3). Gases such as ammonia and carbonyl sulphidehave greater radiative potential but lack efficient formation mechanisms and are photochem-ically unstable in the Martian atmosphere, so they cannot have been present long-term inlarge quantities.Other researchers have looked at hydrogen as a greenhouse gas (Sagan 1977; Ramirezet al. 2014). Wordsworth & Pierrehumbert (2013) showed that N -H CIA can cause sig-nificant warming on terrestrial planets even when H is a minor atmospheric constituent,because broadening of the CIA spectrum at moderate temperatures causes absorption toextend into window regions. In a thought-provoking paper, Ramirez et al. (2014) proposed •• The Climate of Early Mars 9 hat CO -H absorption could have caused warming on early Mars in a similar fashion,perhaps sufficiently to put the climate into a warm and wet state. Hydrogen readily es-capes from a small planet like Mars, so to work this mechanism requires rapid hydrogenoutgassing, which means a very reducing mantle and high rate of volcanism. It also requiresa high atmospheric CO content, which as discussed in Section 3.1 may be inconsistent witha highly reducing mantle (Hirschmann & Withers 2008). Finally, a long-lived highly reduc-ing atmosphere is not obviously consistent with evidence for oxidizing surface conditionsduring the Noachian (Chevrier, Poulet & Bibring 2007). Nonetheless, the contribution ofreducing species to the early Martian climate and atmospheric chemistry is an interestingsubject that requires further research.The radiative forcing of clouds and aerosols was certainly also important to the earlyMartian climate, but it is challenging to constrain. One novel feature of cold CO atmo-spheres is that condensation at high altitudes leads to CO cloud formation (Fig. 5) — aneffect that is observed in the Martian mesosphere today (Montmessin et al. 2007). Forget& Pierrehumbert (1997) proposed that infrared scattering by CO clouds in the high atmo-sphere could have led to significant, long-term greenhouse warming on early Mars. However,to be effective this warming mechanism requires cloud coverage close to 100%. Recent three-dimensional global climate modeling (Forget et al. 2013; Wordsworth et al. 2013) has shownthat this level is never reached in practice. In addition, recent multiple-stream scatteringstudies have indicated that the two-stream methods used previously to calculate CO cloudclimate effects tend to overestimate the strength of the scattering greenhouse (Kitzmann,Patzer & Rauer 2013). Hence the warming effect of CO clouds is likely to have beensmall in reality. Nonetheless, the infrared scattering effect is still an important term intheir overall radiative budget. This means that they at least do not dramatically cool theclimate via shortwave scattering, as was thought to be the case in the earliest studies ofCO condensation on early Mars (Kasting 1991).Recent studies have also investigated the role of H O clouds. In a 3D climate modelstudy, Urata & Toon (2013) found high-altitude water clouds formed that substantiallydecreased the OLR. They proposed that this could have caused transitory or long-livedwarm climate states on early Mars. However, another 3D climate study published in thesame year as Urata & Toon’s work found much less effective upwards transport of watervapor, resulting in mainly low-lying H O clouds that cooled the planet by increasing thealbedo (Wordsworth et al. 2013). It is not unduly surprising that two 3D models of earlyMars produce such different results on cloud forcing, given the uncertainty on this issuefor the present-day Earth (Forster et al. 2007). Nonetheless, the discrepancy highlights theneed for future detailed study of this issue.
The geologic evidence that Mars once had large amounts of surface liquid water is conclusive,but geomorphic constraints on the duration for which that water flowed are much weaker.In addition, much of the geochemical evidence points towards surface conditions that werenot warm and wet for long time periods. This is important, because if repeated transientmelting events are capable of explaining the observations, the theoretical possibilities forwarming become greater.
10 Robin Wordsworth .4.1. Impact-induced Steam Atmospheres.
The late Noachian is coincident with the earlyperiod of intense impactor flux known as the late heavy bombardment (e.g., Hartmann& Neukum 2001). Meteorite impacts were hence unquestionably a major feature of theenvironment during the main period of valley network formation. Based on this, severalproposals for impact-induced warming have been put forward. Segura et al. (2002); Segura,Toon & Colaprete (2008) suggested that large impactors could have evaporated up to tensof meters of water into the atmosphere, which would then have caused erosion when itrained back down to the surface. Later, Segura, McKay & Toon (2012) proposed thatimpact-induced atmospheres could be very long-lived due to a strong decrease in OLR withsurface temperature in a steam atmosphere, which would give rise to a climate bistability.Despite the appealing temporal correlation, impact-induced steam atmospheres are notcompelling as the main explanation for valley network formation. There are two mainreasons for this. First, for transient impact-driven warming, there appears to be a largediscrepancy between the estimated valley network erosion rates (Barnhart, Howard & Moore2009; Hoke, Hynek & Tucker 2011) and the amount of rainfall produced post-impact (Toon,Segura & Zahnle 2010). Second, the runaway greenhouse bistability argument does not seemphysically plausible, at least for clear-sky atmospheres, because it relies on the occurrenceof extreme supersaturation of water vapor in the low atmosphere (Nakajima, Hayashi &Abe 1992). If impacts played a role in carving the valley networks, therefore, they musthave done so by a more indirect method.
Another idea that has seen intensive study over thelast few decades is the SO /H S greenhouse. The martian surface is sulfur-rich (Clark et al.1976) and sulphates are abundant on Hesperian and Noachian terrain (Gendrin et al. 2005;Bibring et al. 2005; Ehlmann & Edwards 2014). This suggests that volcanic emissions ofsulfur-bearing gases (SO and H S) could have had a significant effect on early climate(Postawko & Kuhn 1986; Yung, Nair & Gerstell 1997; Halevy, Zuber & Schrag 2007).SO is a moderately effective greenhouse gas. The 518 cm − (19.3 µ m) vibrational-rotation band associated with its ν bending mode absorbs close the peak of the blackbodyspectrum at 250-300 K, but sufficiently far from the CO ν band at 667 cm − to cause afairly large reduction in the OLR if SO is present at levels of 10 ppm or above (Fig. 3).SO absorption bands in the 1000-1500 cm − region also contribute but partially intersectwith CO CIA at high pressure. H S, which would also have been outgassed in significantquantities by the reducing Martian mantle, is a far less effective greenhouse gas on earlyMars due to the intersection of its main absorption bands with H O and CO (Fig. 3).Like NH and CH , SO is photolyzed in the Martian atmosphere. Photochemicalmodeling under plausible early Martian conditions has suggested that this limits its lifetimeto under a few hundred years (Johnson, Pavlov & Mischna 2009). More importantly, SO photochemistry leads to rapid formation of sulfate aerosols (Tian et al. 2010). These scatterincoming sunlight effectively, raising the albedo and cooling the planet. Similar climateeffects have been observed in stratovolcano eruptions on Earth, such as that of Pinatubo in1991 (Stenchikov et al. 1998).Halevy & Head (2014) argued that intense episodic volcanic SO emissions associatedwith the formation of Hesperian ridged plains could have caused significant greenhousewarming. Using a line-by-line radiative-convective climate model, they found subsolar zonalaverage temperatures of 250 K for SO concentrations of 1-2 ppm in a clear-sky CO -H Oatmosphere. Based on this, they argued that peak daytime equatorial temperatures would ••
Another idea that has seen intensive study over thelast few decades is the SO /H S greenhouse. The martian surface is sulfur-rich (Clark et al.1976) and sulphates are abundant on Hesperian and Noachian terrain (Gendrin et al. 2005;Bibring et al. 2005; Ehlmann & Edwards 2014). This suggests that volcanic emissions ofsulfur-bearing gases (SO and H S) could have had a significant effect on early climate(Postawko & Kuhn 1986; Yung, Nair & Gerstell 1997; Halevy, Zuber & Schrag 2007).SO is a moderately effective greenhouse gas. The 518 cm − (19.3 µ m) vibrational-rotation band associated with its ν bending mode absorbs close the peak of the blackbodyspectrum at 250-300 K, but sufficiently far from the CO ν band at 667 cm − to cause afairly large reduction in the OLR if SO is present at levels of 10 ppm or above (Fig. 3).SO absorption bands in the 1000-1500 cm − region also contribute but partially intersectwith CO CIA at high pressure. H S, which would also have been outgassed in significantquantities by the reducing Martian mantle, is a far less effective greenhouse gas on earlyMars due to the intersection of its main absorption bands with H O and CO (Fig. 3).Like NH and CH , SO is photolyzed in the Martian atmosphere. Photochemicalmodeling under plausible early Martian conditions has suggested that this limits its lifetimeto under a few hundred years (Johnson, Pavlov & Mischna 2009). More importantly, SO photochemistry leads to rapid formation of sulfate aerosols (Tian et al. 2010). These scatterincoming sunlight effectively, raising the albedo and cooling the planet. Similar climateeffects have been observed in stratovolcano eruptions on Earth, such as that of Pinatubo in1991 (Stenchikov et al. 1998).Halevy & Head (2014) argued that intense episodic volcanic SO emissions associatedwith the formation of Hesperian ridged plains could have caused significant greenhousewarming. Using a line-by-line radiative-convective climate model, they found subsolar zonalaverage temperatures of 250 K for SO concentrations of 1-2 ppm in a clear-sky CO -H Oatmosphere. Based on this, they argued that peak daytime equatorial temperatures would •• The Climate of Early Mars 11 xceed 273 K for several months a year during transient pulses of volcanism and hencesignificant meltwater could be generated.Because Halevy & Head (2014) used a one-dimensional column model, they did notaccount for horizontal heat transport by the atmosphere. In contrast to their result, recent3D GCM studies find that in concentrations of under 10 ppm, SO warming cannot causesignificant melting events on early Mars (Mischna et al. 2013; Kerber, Forget & Wordsworth2015). Indeed, the dramatic cooling effects of sulfate aerosols together with the timing ofthe Hesperian flood basalts led Kerber, Forget & Wordsworth (2015) to suggest an oppositeconclusion: the onset of intense sulfur outgassing on Mars may have ended the period ofepisodically or continuously warm conditions in the late Noachian.In summary, no single mechanism is currently accepted as the cause of anomalouswarming events on early Mars. Climate models that allow horizontal temperature variationsshow that in combination, various atmospheric and orbital effects can combine to createmarginally warm conditions and hence small amounts of episodic melting (Richardson &Mischna 2005; Wordsworth et al. 2013; Kite et al. 2013; Mischna et al. 2013; Wordsworthet al. 2015), particularly if the meltwater is assumed to be briny (Fair´en et al. 2009; Fair´en2010). Just like the climate of Earth today, the ancient climate of Mars was probablycomplex, with multiple factors contributing to the mean surface temperature. Nonetheless,further research on this key issue is necessary. The continuing uncertainty with regardto warming mechanisms has recently led some studies to take an empirical approach toconstraining the early Martian climate (Section 4.2).
4. DECIPHERING THE LATE NOACHIAN WATER CYCLE
Radiative-convective climate models are powerful tools, but they have limitations. One ofthe most obvious is their inability to capture cloud effects, except in a crude and highlyparameterized way. A second major limitation is that they fail to account for regionaldifferences in climate, and hence cannot be used to model a planet’s hydrological cycle. Thisis particularly important for Mars, which has large topographic variations and a spatiallyinhomogenous surface record of alteration by liquid water.( While 3D general circulation models (GCMs) of the present-day Martian atmospherebegan to be used from the 1960s onwards (Leovy & Mintz 1969), development of GCMsfor paleoclimate applications has been much slower. An important reason for this is thecomplexity of dense gas CO radiative transfer, as described above. Nonetheless, in the lastfive years, several teams have begun 3D GCM modeling of the early Martian climate.Challenges in simulating the Martian paleoclimate in three dimensions include a poten-tially altered topography compared to present-day, uncertainty in the orbital eccentricityand obliquity, and the difficulty of calculating radiative transfer accurately and rapidly inan atmosphere of poorly constrained composition. Of all of these, the latter poses thegreatest technical challenge. Line-by-line codes such as that used to produce Fig. 3 areimpractical for GCM simulations because the number of calculations required makes themprohibitively expensive computationally. Instead, recent 3D simulations have used thecorrelated- k method (Goody et al. 1989; Lacis & Oinas 1991). This technique, which wasoriginally developed for terrestrial radiative transfer applications, replaces the line-by-lineintegration over spectral wavenumber with a sum over a cumulative probability distributionin a much smaller number of bands.The first study to apply this technique to 3D Martian paleoclimate simulations was
12 Robin Wordsworth ohnson et al. (2008), who looked at the effect of SO warming from volcanism. Whilepioneering in terms of its technical approach, this work used correlated- k coefficients thatwere later found to be in error (Mischna et al. 2013). As a result, it predicted unrealisticallyhigh warming due to SO emissions. Since this time, several new 3D GCM studies of earlyMars using correlated- k radiative transfer have been published, leading to a number of newinsights.As previously discussed, Forget et al. (2013), Wordsworth et al. (2013) and Urata &Toon (2013) investigated the role of CO and H O clouds in warming the early climate,and found both cloud fraction and mean particle size were critical factors in their radiativeeffect. Forget et al. (2013) and Soto et al. (2015) investigated collapse of CO atmospheresdue to surface condensation. They found that the process was extremely significant at lowobliquities (below 20 ◦ ) and at pressures above 3 bar. The predictions of Forget et al. (2013)and Soto et al. (2015) differed in detail, however, partly because Soto et al. (2015) neglectedthe effects of CO clouds and used the present-day solar luminosity. Hence further modelintercomparison on this issue is required. In a related study, Kahre et al. (2013) examinedCO collapse in the presence of an active dust cycle and found that dust could help tostabilize moderately dense atmospheres ( ∼
80 mbar) against collapse at high obliquity.As described in Section 3.4.2, two new 3D studies have also investigated the role of SO warming in the Noachian climate (Mischna et al. 2013; Kerber, Forget & Wordsworth 2015). Another outcome of recent 3D GCM modeling has been an increased understanding of theprocesses governing Mars’ early surface water cycle. On a cold planet, the water cycleis dominated by the transport of surface ice to regions of enhanced stability (cold traps).Ice stability is a strong function of the sublimation rate, which depends exponentially onsurface temperature via the Clausius-Clayperon relation. In practice, this means that theregions of the planet with the lowest annual mean surface temperatures are usually coldtraps. Figuring out the cold trap locations on early Mars is critical, because it tells uswhere the water sources were during warming episodes.Mars today has an atmospheric pressure of around 600 Pa and obliquity of 25.2 ◦ . Themajority of surface and subsurface water ice is found near the poles (Boynton et al. 2002).Mars’ obliquity has varied significantly throughout the Amazonian, however (Laskar &Robutel 1993), and at high obliquities ice migration to equatorial (Mischna et al. 2003;Forget et al. 2006) and mid-latitude (Madeleine et al. 2009) regions is predicted. Obliquityvariations may well have also been important to ice stability in the Noachian and earlyHesperian. However, in this period the role of atmospheric pressure was probably evenmore important.Figure 6 shows the simulated annual mean temperature T s for Mars given a solar flux75% of that today (Gough 1981) and surface pressures of a) p s = 0 .
125 bar and p s = 1 bar.At the lowest pressure, mean surface temperatures are primarily determined by insolationand the main variation of T s is with latitude. At 1 bar, however, a shift to variationof T s with altitude has occurred. The accompanying scatter plot of surface temperaturevs. altitude clearly shows a trend towards temperature-altitude anticorrelation as pressureincreases. The origin of this effect is the increase in coupling between the lower atmosphereand surface via the planetary boundary layer. ••
125 bar and p s = 1 bar.At the lowest pressure, mean surface temperatures are primarily determined by insolationand the main variation of T s is with latitude. At 1 bar, however, a shift to variationof T s with altitude has occurred. The accompanying scatter plot of surface temperaturevs. altitude clearly shows a trend towards temperature-altitude anticorrelation as pressureincreases. The origin of this effect is the increase in coupling between the lower atmosphereand surface via the planetary boundary layer. •• The Climate of Early Mars 13 he surface heat budget on a mainly dry planet can be written as F lw, ↑ = F lw, ↓ + F sw, ↓ + F sens (1)where F lw, ↑ is the upwelling longwave (thermal) radiation from the surface, F lw, ↓ is thedownwelling thermal radiation from the atmosphere to the surface, F sw, ↓ is the incomingsolar flux, and F sens is the sensible heat exchange. The latter term can be written as F sens = ρ a C D c p | u | ( T a − T s ), where ρ a is the atmospheric density near the surface, C D isthe bulk drag coefficient, c p is the atmospheric specific heat capacity at constant pressure, | u | is the surface windspeed and T a is the temperature of the atmosphere at the surface.Observations and simulations indicate that | u | generally decreases with ρ a , but only slowly.Hence the magnitude of F sens will increase significantly with ρ a unless the temperaturedifference T a − T s simultaneously decreases.For a planet without an atmosphere (such as Mercury), F lw, ↓ and F sens equal zero andsurface temperature is determined by solar insolation (with a contribution from geothermalheating in very cold regions). As atmospheric pressure increases, so does heat exchangebetween the atmosphere and surface. For a planet with a thick atmosphere (such as Venusor Earth), sensible and radiative atmospheric heat exchange are significant and drive T s to-wards the local air temperature T a . Because the atmospheric lapse rate follows a convectiveadiabat in the troposphere, surface temperatures therefore decrease with altitude. This isexactly the effect that causes equatorial mountains on Earth such as Kilimanjaro to havesnowy peaks despite the tropical temperatures at their bases.Mars has a lower surface gravity than Earth, which decreases its dry adiabatic lapserate [Γ d = − g/c p , with g gravity and c p the specific heat capacity at constant volume].However, this is more than compensated for by its large topographic variations. The altitudedifference between Hellas Basin and the southern equatorial highlands is around 12 km, orjust over one atmospheric scale height. In 3D GCM simulations, the corresponding drop inannual mean temperature is around 30 K at 1 bar atmospheric pressure, or around 40 K at2 bar (Fig. 6).As a result of this adiabatic cooling effect, for moderate values of Martian obliquity andatmospheric pressures above around 0.5 bar, the equatorial Noachian highlands (where mostvalley networks are observed) become cold traps, confirming a prior prediction by Fastooket al. (2012). Long-term three-dimensional climate simulations coupled to a simple iceevolution model (Wordsworth et al. 2013) have demonstrated that as a result, ice migratesto the valley network source regions regardless of where it is initially located on the surface.The adiabatic cooling effect led Wordsworth et al. to propose the so-called “icy high-lands” scenario for the early climate (Fig. 5). In essence, the idea is that if the valleynetwork source regions were cold traps, the early Martian water cycle could have behavedsomewhat like a transiently forced, overdamped oscillator. Episodic melting events (theperturbing force), would have transported H O to lower altitude regions of the planet onrelatively short timescales. Over longer timescales, adiabatic cooling (the restoring mecha-nism) would have returned the system to equilibrium.Recent modeling and observational work has used the icy highlands hypothesis as aframework for testing various predictions on the early Martian climate. For example, Scan-lon et al. (2013) used an analytical model combined with the 3D GCM of Wordsworth At low atmospheric pressures ( < .
14 Robin Wordsworth t al. (2013) to study orographic precipitation over Warrego Vallis and were able to matchlocal precipitation patterns with the valley network locations. Head & Marchant (2014)discussed the analogies between the icy highlands scenario for early Mars and the AntarcticDry Valleys, which have long been considered one of the most “Mars-like” regions on Earth. (Fastook & Head (2014) used a glacial flow model to study how the buildup of largeice sheets would have affected the geomorphology of the Noachian highlands. Assuming athermal conductivity appropriate to ice without a blanketing effect from snow or firn, theyfound that if the Noachian water inventory was < × the present-day GEL (i.e. 170 m),equatorial highland glaciers would have been cold-based and hence would not have left traceson the surface in the form of cirques, eskers or other glacial landforms. This conclusion wasbroadly confirmed by Cassanelli & Head (2015), who studied the influence of snow and firnthermal blanketing on ice sheet melting under a H O-limited scenario only. The relativelack of glaciation in the equatorial highlands and the total water inventory are key piecesof the Noachian climate puzzle that we return to in Section 4.4.
The icy highlands scenario provides a useful working model for thinking about the earlyMartian climate and fits many aspects of the geologic evidence. Nonetheless, because someobservations are still interpreted as supporting long-term warm and wet conditions, it isalso interesting to model alternative scenarios for the early Martian climate. As we havediscussed, no climate model based on realistic assumptions currently predicts warm andwet conditions for early Mars . To study the hydrological cycle on a warm and wet planet,therefore, it is necessary to use an empirical approach. There are two basic ways to do this:increasing the solar luminosity, or increasing the atmospheric infrared opacity.Wordsworth et al. (2015) tried both these approaches to study precipitation patternson a warm and wet early Mars. As a starting condition, they assumed that Mars hada global northern ocean and smaller bodies of water in Argyre and Hellas, based on theputative delta shoreline constraints of di Achille & Hynek (2010). For simplicity, they alsoneglected the destabilizing effects of the ice-albedo feedback on the early Martian climate(Section 4.3).In general, Wordsworth et al. (2015) found that the precipitation patterns in the warmand wet case were not a close match to the valley network distribution. In particular, thepresence of Tharsis caused a dynamical effect on the circulation that led to low rainfallrates in Margaritifer Sinus, where some of the most well-developed valley networks on Marsare found. This indicates that either a) something was missing from their model or b)Mars was never warm and wet and the Martian valley networks formed through transientmelting events. The opinion of this author is that the second possibility is the correct one.Future work testing the influence of poorly constrained effects such as cloud convectionparametrizations, changes in surface topography and possibly true polar wander on thisresult will allow the first possibility to be tested further. In any case, it is clear thatsystematic empirical investigation of different scenarios using 3D models provides a newway to constrain the early climate. Ramirez et al. (2014) argue that the CO -H CIA mechanism can come close. However, in theirmodel it still requires more atmospheric H than they can produce, even using their most generousoutgassing estimates. ••
The icy highlands scenario provides a useful working model for thinking about the earlyMartian climate and fits many aspects of the geologic evidence. Nonetheless, because someobservations are still interpreted as supporting long-term warm and wet conditions, it isalso interesting to model alternative scenarios for the early Martian climate. As we havediscussed, no climate model based on realistic assumptions currently predicts warm andwet conditions for early Mars . To study the hydrological cycle on a warm and wet planet,therefore, it is necessary to use an empirical approach. There are two basic ways to do this:increasing the solar luminosity, or increasing the atmospheric infrared opacity.Wordsworth et al. (2015) tried both these approaches to study precipitation patternson a warm and wet early Mars. As a starting condition, they assumed that Mars hada global northern ocean and smaller bodies of water in Argyre and Hellas, based on theputative delta shoreline constraints of di Achille & Hynek (2010). For simplicity, they alsoneglected the destabilizing effects of the ice-albedo feedback on the early Martian climate(Section 4.3).In general, Wordsworth et al. (2015) found that the precipitation patterns in the warmand wet case were not a close match to the valley network distribution. In particular, thepresence of Tharsis caused a dynamical effect on the circulation that led to low rainfallrates in Margaritifer Sinus, where some of the most well-developed valley networks on Marsare found. This indicates that either a) something was missing from their model or b)Mars was never warm and wet and the Martian valley networks formed through transientmelting events. The opinion of this author is that the second possibility is the correct one.Future work testing the influence of poorly constrained effects such as cloud convectionparametrizations, changes in surface topography and possibly true polar wander on thisresult will allow the first possibility to be tested further. In any case, it is clear thatsystematic empirical investigation of different scenarios using 3D models provides a newway to constrain the early climate. Ramirez et al. (2014) argue that the CO -H CIA mechanism can come close. However, in theirmodel it still requires more atmospheric H than they can produce, even using their most generousoutgassing estimates. •• The Climate of Early Mars 15 .3. Snowball Mars
An additional impediment to the warm and wet scenario for early Mars that has not yetbeen extensively considered is the ice-albedo feedback (Budyko 1969). This process is animportant player in the ongoing loss of sea ice in the Arctic on Earth due to anthropogenicclimate change (Stroeve et al. 2007). It was also responsible for the snowball Earth globalglaciations that occurred in the Neoproterozoic (Kirschvink 1992; Hoffman et al. 1998;Pierrehumbert et al. 2011).Even with a northern ocean at the di Achille & Hynek (2010) shoreline, Mars wouldpossess a far higher land to ocean ratio than Earth does, making the physics of a snow-ball transition quite different. Unlike on Earth, where runaway glaciation occurs throughfreezing of a near-global ocean, on Mars transport of H O to high altitude regions as snowwould play the key role. Because of the presence of Tharsis at the equator, the topographyof Mars is particularly conducive to an ice-albedo instability of this kind. As discussed inSection 2, most of the formation of Tharsis was probably complete by the late Noachian(see also Fig. 2). Thanks to the adiabatic cooling effect under a thicker atmosphere, Tharsisis an effective cold trap for water ice even if most of Mars is assumed to be warm and wet.To put numbers to this idea, we can define sea level as -2.54 km from the datumfollowing di Achille & Hynek (2010) and take the summit of Tharsis (neglecting the peaksof the Tharsis Montes volcanoes, which are Amazonian-era (Tanaka et al. 2014)) to beapproximately z = 8 km. Then the sea-level-to-summit adiabatic temperature differenceis Γ d z = − gz/c p ≈
38 K. At 1 bar surface pressure the surface temperature gradient withaltitude is still somewhat below the adiabatic value, so we can conservatively estimate atemperature difference of 30 K. To a first approximation, seasonal temperature changes atthe equator can be neglected, so an annual mean sea level temperature of around 30 ◦ C isrequired to avoid the buildup of snow and ice on Tharsis.The high altitude of Tharsis means that the atmospheric radiative effects above it (bothscattering and absorption) are reduced. This increases its radiative forcing due to surfacealbedo changes compared to lower-lying regions. Because of its equatorial location, anice-covered Tharsis also increases the planetary albedo regardless of Mars’ obliquity.No 3D climate simulation has yet addressed the impact of a snowy Tharsis on thestability of a warm and wet early Mars. This is an important topic for future study. Itmay be that marginally warm global mean temperatures are insufficient to stop Mars fromrapidly transitioning to a cold and wet state.
As discussed in Section 2, there is little evidence for glaciation during Mars’ Noachian periodin the equatorial highlands. At first glance, this seems like a potential drawback of the icyhighlands scenario for the early climate. If snow and ice buildup in high altitude regionswas the source of the water that carved the valley networks, should we not see a morewidespread record of glacial and fluvioglacial alteration across the Noachian highlands?The trouble with this line of thinking as an argument against the icy highlands scenariois that it equally applies to a warm and wet early Mars. If Mars was once warm andwet, it must have subsequently become cold, because it is cold today. Once it did, liquidwater would freeze into ice and migrate to the cold trap regions of the surface. Thisimplies the buildup of ice sheets in the equatorial highlands unless the Martian atmosphericpressure immediately decreased to low values and obliquity was continually low in the period
16 Robin Wordsworth mmediately following the warm and wet phase (which seems unlikely). The quantity ofsurface water sufficient to fill a northern ocean to the di Achille & Hynek (2010) shorelineimplies a GEL of around 550 m, so in the immediate post-warm-and-wet phase these icesheets could be several kilometers thick. As shown by Fastook & Head (2014), this wouldlead to wet-based glaciation (and hence fluvioglacial erosion) even under very climate coldconditions, and even without the insulating effects of a surface snow layer.If a warm and wet Mars should leave abundant evidence of glaciation in its subsequentcold and wet phase, how can the equatorial periglacial paradox be explained? Most likely,the resolution lies in early Mars’ total surface H O inventory. In a supply-limited icyhighlands scenario with episodic melting, snow and ice deposits are cold-based. Then, theonly significant alteration of terrain comes from fluvial erosion during melting events.Many studies have investigated the evolution of Mars’ surface H O inventory throughtime. The deuterium to hydrogen (D/H) ratio can be used to constrain the early Martianwater inventory (Greenwood et al. 2008; Webster et al. 2013; Villanueva et al. 2015), al-though our lack of knowledge of the dominant escape process in the Noachian/Hesperianand of the cometary contribution to Mars’ surface water (Marty 2012) complicates theanalysis. Villanueva et al. (2015) recently used Earth-based observations of Martian D/Hand estimated the late Noachian water GEL to be 137 m. In contrast, Carr & Head (2015)recently calculated the H O loss/gain budget in the Hesperian and Amazonian and con-cluded that the late Noachian water GEL was as low as 24 m. Although the uncertainty isconsiderable, most estimates place the early Martian water inventory below a few hundredmeters GEL. This low inventory compared to Earth’s is consistent with Mars’ low mass andlikely significant loss of atmosphere to space (see sidebar).Continuing our previous line of thought to its logical conclusion, we can constructan idealized two dimensional phase diagram for the early Martian climate, with surfacetemperature on one axis and the total surface H O inventory on the other (Figure 7).Each quadrant in Fig. 7 represents a different end-member state for the long-term climateand surface hydrology of early Mars. The warm and wet state with a northern ocean isdisfavored for the geochemical and climatological reasons discussed previously. Becauseof the ice-albedo feedback, it also readily transitions to a cold and wet state. The coldand wet state implies extensive wet-based glaciation across Noachian terrain, in conflictwith the geomorphological record. The cold and (relatively) dry state (water GEL <
200 m)is essentially the icy highlands scenario, which in combination with the right amount ofepisodic melting can explain most of the geologic record. Finally, there is a possible “warmand dry” scenario where all the surface water is liquid but the H O GEL is below ∼
200 mis also possible. Potentially, this might also fit many of the geologic constraints on the earlyMartian climate. However, it is climatologically as hard to justify as the warm and wetstate. In addition, in such a scenario the low-lying regions where liquid water stabilizeswould be so far from the valley network source regions that precipitation there might belimited or non-existent. Preliminary GCM studies using the model described in Wordsworthet al. (2015) suggest that this is indeed the case. This is yet another important issue thatdeserves to be studied in detail in future.
5. OUTLOOK
While major questions remain on the nature of the early Martian climate, recent advanceshave been significant. The weight of geomorphological and geochemical evidence points ••
While major questions remain on the nature of the early Martian climate, recent advanceshave been significant. The weight of geomorphological and geochemical evidence points •• The Climate of Early Mars 17
ARS, THE RUNT PLANET
Mars is fascinating to study in its own right, but also because because of the insight it can give us into plan-etary evolution and habitability in general. Several potentially rocky exoplanets and exoplanet candidatesthat receive approximately the same stellar flux as Mars have now been discovered [e.g., Udry et al. (2007);Quintana et al. (2014)]. Can we derive lessons from Martian climate evolution studies that are generalizableto a wider context?Mars formed very rapidly compared to Earth (Dauphas & Pourmand 2011) and is smaller than pre-dicted by standard formation models (Chambers & Wetherill 1998). Current thinking suggests Mars is bestunderstood as a planetary embryo, whose development into a fully fledged terrestrial planet was arrestedby some process such as instability in the configuration of the outer planets (Walsh et al. 2011). The RedPlanet is not as massive as it could have been, and this factor more than any other has dominated itssubsequent evolution.Mars’ low mass relative to Earth led to an early shutdown of the magnetic dynamo (Acuna et al.1998) and rapid cooling of the interior. If plate tectonics ever initiated at all, it ceased rapidly (Connerneyet al. 2001; Solomon et al. 2005). As a result, volatile cycling between the surface and interior was stronglyinhibited and the rate of atmospheric loss to space was probably also enhanced. The present-day atmosphereis so thin that liquid water is unstable on the surface. Mars’ reddish, hyperoxidised surface, which isextremely inhospitable to life, is a direct result of the escape of hydrogen to space over geologic time(Lammer et al. 2003).Almost certainly, Mars tells us more about the habitability of low mass planets than the habitability ofplanets that are far from their host stars. Indeed, an Earth-mass planet with plate tectonics at Mars’ orbitaldistance would be habitable today, if mainstream thinking on the carbonate-silicate cycle (Walker, Hayes& Kasting 1981) and planetary habitability (Kasting, Whitmire & Reynolds 1993) is correct. Radiative-convective calculations using the model described in Wordsworth & Pierrehumbert (2013) indicate thata 1 M ⊕ planet at Mars’ orbit with an atmosphere dominated by CO and H O would have a surfacetemperature of 288 K for an atmospheric pressure of around 3-5 bar – a small amount in comparison withEarth’s total carbon inventory (Sleep & Zahnle 2001; Hayes & Waldbauer 2006).In the absence of the still-mysterious process that abruptly halted Mars’ growth, our more distantneighbour could have been globally habitable to microbial life through much of its history. Given thepotential for exchange of biological material on impact ejecta between terrestrial planets (Mileikowsky et al.2000), biogenesis on Earth [or Mars; Kirschvink & Weiss (2002)] would then have led to the developmentof global biospheres on two planets in the Solar System. A future scenario in which Mars possesses a globalbiosphere is also possible, but will depend on human colonization and subsequent intentional modificationof the climate. towards a late Noachian hydrological cycle that was intermittent, not permanently active.Three-dimensional GCM simulations of the early climate and other modeling and analoguestudies suggest a water-limited early Mars with episodic melting episodes may be a suitableparadigm for much of the late Noachian and early Hesperian climate.Despite the progress over the last few years, aspects of the early climate remain un-clear. Chief among these is the driving mechanism behind the episodic warming events.Previous theories for warming by CO clouds, sulfur volcanism and impact-induced steamatmospheres have all been shown to possess serious problems. Indeed, as one- and three-
18 Robin Wordsworth imensional climate models become more accurate, there has been a general trend towardsprediction of colder surface temperatures. A successful theory for surface warming mustyield valley network erosion rates consistent with the geologic evidence (Hoke, Hynek &Tucker 2011) but avoid extensive surface carbonate formation and aqueous alteration ofsediments (Tosca & Knoll 2009; Ehlmann et al. 2011).The climate of early Mars, like that of present-day Earth, almost certainly involvedmultiple interacting processes. In addition, despite the difficulties with the impact-inducedsteam atmosphere hypothesis, the timing of the peak period of valley network formationwith respect to the late heavy bombardment is suggestive of a causal link. In contrast,the lack of temporal correlation between the valley networks and Hesperian ridged plainsargues against a causal link with effusive volcanism.The key advantage of modeling the early climate as a spatially varying system, ratherthan studying isolated processes or trying to achieve mean temperatures above 273 K in aglobally averaged model is that it allows tighter intercomparison with the geologic evidence.Future research should work towards close integration of surface geology, 3D climate andhydrological modeling studies. Specific regions where this approach would be particularlyuseful include the south pole around the Dorsa Argentea Formation and the Aeolis quadran-gle where Gale crater is located. Issues of specific importance to the 3D climate modelinginclude better representation of clouds, perhaps through a sub-gridscale parametrizationscheme [e.g., Khairoutdinov & Randall (2001)] and coupling with subsurface hydrologymodels (Clifford 1993; Clifford & Parker 2001).If Mars never had a steady-state warm and wet climate, does this spell doom for thesearch for past life on the surface? Probably not: life on Earth clings tenaciously to almostany environment where we can look for it, including Antarctica’s Dry Valleys and deepbelow the seafloor (Cary et al. 2010; D’Hondt et al. 2004). Indeed, if one view of Earth’sclimate in the Hadean and early Archean is correct, our own planet may have been in aglobally glaciated state when life first formed (Sleep & Zahnle 2001). The search for life onMars must continue, but to maximize the chances of success it needs to be informed by ourevolving understanding of the early climate. (
SUMMARY POINTS
1. Mars underwent an extended period of surface erosion and chemical weathering byliquid water until around 3.5 Ga, during the late Noachian and early Hesperianperiods.2. The weight of the observational evidence favors a mainly cold climate with episodicwarming events, rather than permanently warm and wet conditions.3. If early Mars was once warm and wet, thick wet-based icesheets would have formedon the Noachian highlands when the warm period ended, causing significant glacialerosion.4. Constraints on the early solar luminosity, Martian orbit and radiative transfer ofCO strongly disfavor a warmer climate due to CO and H O only.5. Under a thicker atmosphere, adiabatic cooling of the surface causes transport ofsnow and ice to the valley network source regions.6. Repeated episodic warming events probably caused ice and snowpacks in theNoachian highlands to melt, carving valley networks and other fluvial features.7. The precise mechanism that caused the warming events is still poorly constrained. ••
1. Mars underwent an extended period of surface erosion and chemical weathering byliquid water until around 3.5 Ga, during the late Noachian and early Hesperianperiods.2. The weight of the observational evidence favors a mainly cold climate with episodicwarming events, rather than permanently warm and wet conditions.3. If early Mars was once warm and wet, thick wet-based icesheets would have formedon the Noachian highlands when the warm period ended, causing significant glacialerosion.4. Constraints on the early solar luminosity, Martian orbit and radiative transfer ofCO strongly disfavor a warmer climate due to CO and H O only.5. Under a thicker atmosphere, adiabatic cooling of the surface causes transport ofsnow and ice to the valley network source regions.6. Repeated episodic warming events probably caused ice and snowpacks in theNoachian highlands to melt, carving valley networks and other fluvial features.7. The precise mechanism that caused the warming events is still poorly constrained. •• The Climate of Early Mars 19 i g u r e ( l e f t) C o n t o u r p l o t o f t h e p r e s e n t - d a y M a r t i a n t o p og r a ph y f r o m M O L A d a t a . M a j o r f e a t u r e s o f t h e t e rr a i n a r e i nd i c a t e d . ( b o tt o m l e f t) Sp a t i a l d i s t r i bu t i o n o f t h e t h r ee m a i n t e rr a i n t y p e s o n t h e M a r t i a n s u r f a ce [ d a t a f r o m T a n a k a ( ) ; S c o tt & T a n a k a ( ) ; T a n a k a & S c o tt( ) ]. ( r i g h t) H i g h li g h t s o f t h e g e o m o r ph o l og i c a l e v i d e n ce f o r a n a l t e r e d c li m a t e i n t h e N oa c h i a n a nd H e s p e r i a n . ( A ) F l u v i a l c o n g l o m e r a t e s o b s e r v e d i n s i t u b y t h e C u r i o s i t y r o v e r a t G a l e C r a t e r [ f r o m W illi a m s e t a l. ( ) ; r e p r i n t e d w i t hp e r m i ss i o n f r o m AAA S ]. ( B ) D e l t a i c l a k e d e p o s i t s i n E b e r s w a l d ec r a t e r [ f r o m M a li n & E d g e tt( ) ; r e p r i n t e d w i t hp e r m i ss i o n f r o m AAA S ]. ( C ) V a ll e y n e t w o r k s i n P a r a n ´a V a lli s [ f r o m H o w a r d , M oo r e & I r w i n ( ) ]. ( D ) S i nu o u s r i d g e s i n t h e D o r s a A r g e n t e a F o r m a t i o n i n t e r p r e t e d a s g l a c i a l e s k e r s [ f r o m H e a d & P r a tt( ) ].
20 Robin Wordsworth
RE- NOACHIAN
NOACHIAN HESPERIAN AMAZONIAN HADEAN ARCHEAN PROTEROZOIC PHANEROZOIC Volcanism
Mars Earth
Impactors Fluvial/Glacial Surface AlteraDon Aqueous Mineralogy ² HELLAS (DEFINES START OF NOACHIAN PERIOD) ² ISIDIS ² ARGYRE VALLEY NETWORKS CRUST FORMATION Fe/Mg PHYLLOSILICATES BOREALIS?
THARSIS FORMATION HESPERIAN RIDGED PLAINS OUTFLOW CHANNELS DORSA ARGENTEA FORMATION SULFATES AND EVAPORITES ELEVATED IMPACTOR FLUX ? ? ? Al PHYLLOSILICATES
Figure 2
Timeline of major events in Mars history, with the geologic eons of Earth displayed above. Ingeneral, the absolute timing of events on Mars is subject to considerable uncertainty, but thesequencing is much more robust. Question marks indicate cases where processes could also haveoccurred earlier but the geologic record is obscured by subsequent events. Based on data fromWerner & Tanaka (2011); Fassett & Head (2011); Ehlmann et al. (2011) and Head & Pratt (2001).
The two most likely forcing mechanisms are meteorite impacts and volcanism, al-though the details remain unclear.
DISCLOSURE STATEMENT
The author is not aware of any affiliations, memberships, funding, or financial holdings thatmight be perceived as affecting the objectivity of this review.
ACKNOWLEDGMENTS
The author thanks Raymond Pierrehumbert for a constructive review of this manuscript,and numerous colleagues for their critical feedback and advice on key aspects of the obser-vations, including Bethany Ehlmann, Caleb Fassett, Jim Head, Francois Forget and LauraKerber. Bob Haberle is also acknowledged for enlightening discussion of the “warm anddry” state for the early climate. ••
The author thanks Raymond Pierrehumbert for a constructive review of this manuscript,and numerous colleagues for their critical feedback and advice on key aspects of the obser-vations, including Bethany Ehlmann, Caleb Fassett, Jim Head, Francois Forget and LauraKerber. Bob Haberle is also acknowledged for enlightening discussion of the “warm anddry” state for the early climate. •• The Climate of Early Mars 21 igure 3
The greenhouse effect on early Mars. (top) Absorption cross-sections per molecule of backgroundgas vs. wavenumber at 1 bar and 250 K, for various greenhouse gases in the early Martianatmosphere, with the gas abundances given in the legend. Results were produced using theopen-source software kspectrum . (bottom) OLR vs. wavenumber from early Mars calculated usinga line-by-line calculation assuming surface pressure of 1 bar, surface temperature of 250 K and a167 K isothermal stratosphere. Blackbody emission at 250 K and 167 K is indicated by thedashed lines. The black line shows OLR for a pure CO atmosphere, while the red line showsOLR with all the additional greenhouse gases in the top plot included. Results were producedusing the HITRAN 2012 database, the Clough, Iacono & Moncet (1992) approach to solving theinfrared radiative transfer equation and the GBB CO CIA parametrization from Wordsworth,Forget & Eymet (2010).
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Icarus ••