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Question 10: Dynamic Habitability

The past several decades of exploration of Earth’s biosphere have expanded our knowledge of the range of environments with liquid water, nutrients, and energy sources that sustain life.¹ Simultaneously, the past several decades of planetary exploration have revealed multiple ancient and modern potential habitable environments across the solar system. The study of planetary habitability requires understanding the factors controlling habitability, the evolutionary pathways by which an environment became habitable, and the processes that sustain, enhance, diminish, or even extinguish habitability. Increasingly, we understand that habitability is not a yes/no proposition but a continuum—that is, more habitable, less habitable, not habitable—and that planetary environments transition from habitable to not habitable, and vice versa, over space and time. Feedbacks and interrelationships between stellar, dynamical, and planetary evolution drive environmental change and changes in habitability. Life itself can also modify environments and their habitability (Question 9, Chapter 12). Understanding the dynamic habitability of bodies in the solar system through study of their past, current, and future evolutionary trajectories lays the foundation for the study of potentially habitable worlds beyond the solar system.

Here, the committee reviews the state of knowledge of habitability, including habitability of solar system worlds. The committee then examines key aspects controlling habitability: the availability of water, organics, nutrients (e.g., the main chemical elements, CHNOPS), and energy as well as the role of time and the stability or continuity of habitability—that is, the dynamical component of habitability. The chapter concludes with recommended activities and measurements to advance knowledge of solar system habitability.

Q10.1 WHAT IS HABITABILITY?

Planetary habitability is the measure of a body’s potential to develop and sustain life. Because there is no current example of life beyond Earth, planetary habitability is largely an extrapolation of conditions under which we find life on Earth (Question 9), as well as knowledge of the characteristics of the Sun and solar system that appear favorable for life to develop and flourish.

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¹ A glossary of acronyms and technical terms can be found in Appendix F.

At present, with a single example of an inhabited world (Earth), we lack the information to fully understand the conditions that lead to—or prohibit—the origin and sustenance of life on worlds in the solar system. Through our study of life on Earth, however, we can measure the conditions under which life as we know it survives, thus providing a benchmark for studying habitability. The search for current or past habitable conditions on other solar system bodies is thus essential to understand both the fundamental processes governing habitability as well as whether life exists on worlds beyond our own.

Q10.1a What Are the Environmental Characteristics Required for Habitability?

The principal habitability criteria at the planetary scale are “the presence of liquid water, conditions favorable for the assembly of complex organic molecules at some time during the planet’s history, and energy sources to sustain metabolism” (NASEM 2019a). Favorable conditions include both the materials necessary for life and the maintenance of environmental conditions conducive to life over time (Figure 13-1). The concept of dynamic habitability extends this framework, incorporating the exogenic and endogenic processes that change over time to generate and sustain—or not—the principal habitability criteria. At the planetary scale, dynamic habitability includes “stellar evolution and its impact on the presence of liquid water over time; the evolving structure of the planet interior influencing the magnetic field and plate tectonics; and the state of the atmosphere, which might, for example, redistribute metabolic energy” (NASEM 2019a) (Figure 13-2).

A traditional definition of the habitable zone of a planetary system centers on surface liquid water: the radial distance from a star where incident irradiance, potentially aided by an atmospheric greenhouse, results in temperatures that (given sufficient atmospheric pressure) permit liquid water to remain stable on a planetary surface. The stability of habitable surface environments depends on feedbacks between stellar, orbital, atmospheric, and geological evolution (see Q10.7). Although this definition is useful to identify potentially habitable worlds by remote sensing (a necessity for exoplanets), habitable worlds need not have stable surface liquid water at the surface. Subsurface habitable environments exist on Earth and subsurface liquid water exists on other solar system bodies, mostly independently of solar heating, with water kept liquid by geothermal heating provided from accretion, impacts, radioactivity, volcanism,

or tidal dissipation, and with geochemical sources of energy for metabolic activity. On rocky bodies, liquid water can be hosted in regolith and in the deeper subsurface. On icy bodies, liquid water can exist in subsurface oceans that are directly in contact with a rocky core or sandwiched between two ice layers, or in reservoirs within the outer ice shell. Liquid water also exists as droplets within atmospheres and, perhaps, within the subsurface of large asteroids (Q10.3). As of 2020, several dozen exoplanets have been identified within the habitable zone of their parent star, where surface temperatures may allow liquid water to be present (Question 12, Chapter 15). Many more may harbor liquid water in their subsurface, but this property would be difficult to detect remotely (Quick et al. 2020).

The habitability of these aqueous environments depends on the supply (abundance and fluxes) of organic matter, chemical nutrients, and energy (Q10.4–10.6), as well as on their longevity (reviewed in Q10.7). Crucially, the exploration of the past, current, and future habitability of a world is distinct from the question of whether that world has actually harbored life. Understanding the factors that generate habitability and how often a habitable environment is or can become inhabited is central to understanding the prospect for how prevalent and distributed life is in the universe.

Q10.1b How Is Habitability Sustained, Changed, or Lost Over Time?

The solar system is rife with examples of changing habitability over time. The sustainability of habitability on terrestrial planets is determined by planetary size, starting composition, interior evolution (including, but not limited to, dynamo generation), impact bombardment, changes in stellar flux, and orbital dynamics (see Figure 13-2).

For example, any early ocean and atmosphere on Earth would have been sterilized by the Moon-forming impact. Subsequently, as Earth cooled and water oceans became reestablished, early Earth may have had a CO₂- or a CH₄-rich atmosphere akin to Titan today, depending on the balance of input from volcanism and impacts. One of the biggest atmospheric changes on Earth, the rise of oxygen ~2.4–2.0 Ga, was biologically driven. (Additionally, Earth’s geologic record indicates periods where much larger portions of Earth’s land and ocean were ice-covered during profound glaciations, coined “Snowball Earth” episodes.) On Mars and Venus, substantial changes in habitability are expressed in atmospheric isotopes and geological records. Mars’s current 6-mbar CO₂ atmosphere does not sustain stable surface liquid water, but water-deposited sediments, hydrous minerals, and fluvial and glacial landforms visible from orbit attest to the past presence of lakes, rivers, aquifers, and hydrothermal systems, as well as past glaciations. Venus’s current dense CO₂ atmosphere sustains a high-temperature greenhouse effect at the surface where liquid water cannot exist, but atmospheric isotopes and some evolutionary models hint at a more clement past (Q6.2b). Oceans have been confirmed on jovian and saturnian satellites, although it is not understood whether, for example, the ocean on tiny Enceladus has the same prospect for long-term heating as larger Europa. On icy worlds, the extent to which accretional, differentiation, radiogenic, and tidal heating can sustain subsurface oceans supplied with chemical energy and nutrients varies over time and can reflect the coevolution of satellites’ orbits and interiors in a circumplanetary system (Question 8, Chapter 11).

Strategic Research for Q10.1

• Characterize conditions on known and candidate modern habitable environments in the solar system—for example, on Mars, Enceladus, Europa, and other bodies—by measuring water chemistry, mineralogy, ice composition, gases, and organic molecules to assess whether conditions exist that could support life.
• Determine whether there are modern habitable environments in atmospheres by characterizing chemistry, including organic molecules, in the atmospheres of Venus and Titan.
• Determine the character, timing, and duration of past habitable environments on Mars using chemical, mineralogical, textural, isotopic, and organic measurements from orbit, in situ, and on returned samples.
• Determine whether Venus ever hosted liquid water on its surface by geomorphic mapping to search for water-formed landforms as well as mineralogy, chemistry, and isotopic measurements in situ or with samples that may record crust interaction with water and volatile evolution over time.
• Understand interior structures, tidal dissipation dynamics, and surface-interior exchange for icy shells of ocean worlds via measurement by spacecraft, theory, and modeling to determine the magnitudes and timescales of heating and persistence of liquid water.

Q10.2 WHERE ARE OR WERE THE SOLAR SYSTEM’S PAST OR PRESENT HABITABLE ENVIRONMENTS?

The solar system offers a broad range of locales to investigate the dynamic nature of habitable environments (Table 13-1): planetary bodies recognized as habitable today, those that were perhaps once habitable, and those that provide information on the framework of habitability. Below, the committee identifies priorities for developing a deeper understanding of pathways toward habitability.

Q10.2a What Can Terrestrial Planetary Bodies Reveal About Habitable Environments?

Rocky worlds in the inner solar system provide key context for understanding interior–surface–atmosphere interactions, including the role played by water in contributing to planetary habitability.

Mercury and the Moon

Although Mercury and the Moon do not possess habitable environments, they provide insights on the processes controlling habitability. Both host ice deposits in the permanently shadowed regions of their poles. MESSENGER revealed that Mercury’s poles also host organic-rich deposits that can inform sources of organics on inner planets.

TABLE 13-1 Question 10. The Factors That Govern Planetary Habitability, and Whether Those Factors Are Present for Select Planetary Bodies Across the Solar System

NOTE: Cells with a check and question mark signify likely/probable.
SOURCE: Courtesy of P.K. Byrne.

Both airless worlds can help us understand the role of volcanism in the accumulation and nature of an atmosphere, and how such an atmosphere can be lost. Their cratering records establish a flux-calibrated chronology (Question 4, Chapter 7) informing how changing impact bombardment has influenced habitable conditions, especially in the early inner solar system when larger impacts were more frequent. Mercury’s current intrinsic magnetic field enables the study of how interior evolution sustains such fields over time, as well as how a magnetic field moderates physical and chemical interactions between the surface and the space environment.

Venus

Venus holds critical clues for understanding the role of climate change for a world becoming uninhabitable. Venus’s 740 K surface is clearly uninhabitable today. Measurements of its atmosphere by previous missions, including most recently Venus Express, have suggested that Venus lost perhaps as much as an ocean on Earth’s worth of water over its lifetime. When, how, and how much water was lost remain key open questions. One possibility is that shortly after formation, accretional and impact-generated heat, aided by close proximity to the Sun, led Venus to have a magma ocean that outgassed a dense steam atmosphere (e.g., Kasting 1988; Hamano et al. 2013). Under this scenario, Venus’s surface was never habitable. However, if Venus escaped the magma ocean phase without experiencing such a runaway greenhouse effect, a cloudy subsolar hemisphere and slow rotation could have helped keep the planet cool and clement for potentially billions of years (Way and Del Genio 2020). In this model, Venus only underwent runaway greenhouse warming within perhaps the past billion years, triggered by a massive release of CO₂ into the atmosphere by volcanic eruptions. If so, there may have been two Earth-sized habitable worlds for much of solar system history, as seen in some

exoplanetary systems (Question 12, Chapter 15). Whether Venus’s surface was once habitable, and when it ceased to be, could be determined by studying atmospheric isotopes and the mineralogy, chemistry, and isotopes of the rock record.

Even today, Venus’s atmosphere has been suggested as a prospective abode for life; at altitudes of 55–60 km, the atmospheric pressure and temperature are comparable to Earth at sea level. The presence of a time-variable ultraviolet-absorbing material and unexplained aspects of measured gas chemistry have led to speculation about a role for biology (Q9.1c). Chemical reactions in a very acidic, low-water, high ultraviolet atmosphere and potential for its habitability needs to be quantified by in situ compositional measurements.

Mars

Mars offers a basis for understanding rock-water-atmosphere interactions and long-term climate change on a body that hosted habitable environments for the first 1–2 billion years of its history and may still host habitable environments. Morphological and mineralogical orbital data reveal a rich record of liquid water with a variety of habitable environments: past hydrothermal springs, groundwater aquifers, deep open and closed basin lakes, sulfate and chloride playa lakes, weathered soils, and possibly oceans, varying geographically and temporally (e.g., Ehlmann and Edwards 2014). In situ results from landed missions have investigated a few of these past environments in detail. The Opportunity rover found evidence for past sulfate playa lakes with acidic waters, which on Earth can be inhabited by organisms capable of tolerating low pH. The Spirit rover identified silica deposits interpreted to have formed in an ancient volcanic hydrothermal spring; such environments are populated with microbial life on Earth. The Curiosity rover found evidence for a ~3.5 Ga habitable lake system within Gale crater with low salinity, neutral pH, and all chemical elements essential to life, as well as episodic pulses of groundwater, until at least ~2.7 Ga. Curiosity also identified long-chain organic material within the sedimentary rocks (e.g., Eigenbrode et al. 2018; Rampe et al. 2020). What processes sustained these habitats with what continuity, and why did conditions change? Surface water in the past was supported by a thicker and warmer atmosphere, since lost. Results from the MAVEN orbiter indicate loss of the martian atmosphere to space (Jakosky et al. 2018), and chemical weathering was a major sink for water and other volatiles (Scheller et al. 2021). The habitability of early Mars motivates the ongoing search for biosignatures within its multiple, diverse ancient aqueous environments (Question 11, Chapter 14). The Mars Sample Return campaign will allow for martian samples acquired by the Perseverance rover to be returned to Earth laboratories for detailed analyses that address the open questions about what processes supported the formation of habitats in Mars’s past.

Today, Mars is cold and arid, and whether it hosts modern habitable environments is a key question. Unlike Earth, Mars has large, obliquity (spin axis tilt) variations that drive substantial climate excursions on hundred-thousand-year timescales. These excursions periodically release volatiles into the atmosphere, raising atmospheric temperature and pressure above what we observe today and potentially making ephemeral occurrence of liquid water possible. Orbiters have identified ice sheets and glaciers distributed from the poles to the mid-latitudes and ground ice within 1 m of the surface (Morgan et al. 2021). Whether liquid water exists beneath such ice is not yet known. Features such as recurring slope lineae and other downslope debris flows that preferentially occur during warm months have characteristics suggesting the involvement of liquid water, even if a purely dry process is considered the most likely explanation (McEwen et al. 2021). The martian regolith contains salts, some geologically young (Wray 2020), which may melt ice (Stillman and Grimm 2011) or liquefy in the presence of atmospheric water vapor (i.e., deliquescence) (Rivera-Valentín et al. 2020). At Gale crater, continued in situ detection and cyclical seasonal variation of parts-per-billion traces of methane (Webster et al. 2018)—a nonequilibrium gas that can provide a source of energy for life and, on Earth, can be produced by life—contrasts with its nondetection at parts-per-trillion level from orbit (Korablev et al. 2019). A localized source and differences in time of day of measurement may reconcile these measurements (Moores et al. 2019; Webster et al. 2021), but the process generating methane is not yet known. Collectively, these observations sustain the debate on the habitability of present-day Mars and whether it could be inhabited.

Q10.2b Which Icy Ocean Worlds and Dwarf Planets Are Habitable?

Missions to the outer solar system have demonstrated the presence of organic chemistry and liquid water currently or in the past under the surfaces of many objects, including in subsurface oceans within icy moons

around the giant planets. The habitable zone, as defined by the presence of liquid water (Q10.1), thus extends considerably farther out in the solar system than suggested by insolation limits on surface water (Coustenis and Encrenaz 2013).

Although key questions concerning their longevity and habitability remain outstanding, identifying the distribution of subsurface oceans is an important first step. Confirmed or strongly suspected ocean worlds, determined primarily from Galileo and Cassini-Huygens measurements, include (Lunine 2017; Hendrix et al. 2019; Hand et al. 2020):

• Enceladus: Plumes of water from fractures at this moon’s south pole escape from a global liquid water ocean, bearing silica, salts, and organics that indicate chemical reactions between that ocean and a silicate interior (Postberg et al. 2018a). Key ingredients for habitability appear to be met for this body (Figure 13-1). A porous, rocky interior may sustain these environments through time. Key remaining questions concern the longevity of the ocean and its variability with time.
• Europa: Observations of a young and smooth surface with crisscrossing fractures and magnetic induction measurements together provide strong evidence for a global ocean beneath a geologically active ice shell (e.g., Kivelson et al. 2009). Key questions concern the composition of the ocean, the nature of its interactions with the silicate interior, the presence and inventory of organic compounds and bioessential elements, and how the ocean undergoes physical and geochemical exchange with the overlying ice-shell.
• Ganymede and Callisto: Magnetic induction measurements on Ganymede revealed a global ocean, likely sandwiched between two or more ice layers with deep, high-pressure ice phases. Ganymede also has an intrinsic magnetic field. Although its precise internal structure is not known today, an induction response in Callisto has been suggested, but may be due either to an ocean or to ionospheric interference (e.g., Hartkorn and Saur 2017).
• Titan: A global deep ocean is thought to be present within Titan between ice layers (e.g., Iess et al. 2012). The surface of Titan’s ice shell is covered with deposits of liquid hydrocarbons. The nitrogen–methane atmosphere has a rich methanological cycle and supports complex organic chemistry that starts in the ionosphere and diffuses down to the surface, forming prebiotic molecules and aerosols/hazes (e.g., Coustenis 2021). Cryovolcanism has been hypothesized and could deliver methane from the interior to the atmosphere (Lopes et al. 2013).

More putative ocean worlds await evidence to conclusively confirm suggestions from modeling and observations of having had subsurface oceans. Beyond liquid water oceans, findings over the past decade have revealed a continuum that may include “mudballs”—that is, porous, rocky interiors filled with silicate-brines, which increase water–rock interfaces at which energy and nutrients may be supplied. Candidate ocean worlds include:

• Triton: The geologically young surface of Triton has been linked to convective and/or cryovolcanic processes, consistent with the observation by Voyager 2 of active nitrogen plumes and, potentially, a subsurface ocean (e.g., Hansen et al. 2021).
• Ceres: With a pervasively hydrous and icy crust, and evidence for modern-day brines or diapirs generating young salty deposits on its surface, Ceres once had, and may still have, subsurface liquid water (e.g., Castillo-Rogez et al. 2020). Ceres appears to have an abundant supply of organic compounds and nitrogen at the surface, but its energy supply may be scarce as it is not subject to tidal heating.
• Uranian and other saturnian moons (especially Ariel and Miranda, or Dione): These moons have dynamic geological histories preserved on their surfaces, and modeling efforts suggest the possibility that liquid water or mud oceans may exist in the subsurface (e.g., Bierson and Nimmo 2022). Carbon and nitrogen have been detected at Ariel’s surface.
• Pluto: There is evidence that Pluto currently possesses a liquid water ocean beneath its thick frozen ice shell, which developed when ice melted owing to heat from radioactive elements in Pluto’s core, long after the body itself formed (e.g., Nimmo and McKinnon 2021). Pluto’s surface and atmosphere are also rich in carbon and nitrogen.

Q10.2c What Can Small Bodies Reveal About Habitability, and Are There Habitable Dwarf Planets Today?

Results over the past decade by NASA’s Dawn and New Horizons missions revolutionized our understanding of habitability related to small solar system bodies, such as asteroids, comets, and trans-neptunian objects (TNOs). In addition to potential subsurface water at Ceres and Pluto, ground-based observations of dwarf planets hint at the potential for more ocean worlds in the trans-neptunian population. Determining the occurrence and characteristics of liquid water environments in this population of small bodies remains an open area of investigation.

Additionally, the meteorite record suggests that some asteroids may have experienced aqueous alteration, which could have supported prebiotic chemistry. The presence of water on some asteroids was confirmed this past decade (Campins et al. 2010). Given their observed comet-like activity (e.g., dust tails and other mass loss events), some active asteroids may also presently possess water ice. Although small bodies themselves are unlikely to sustain life, the chemical reactions facilitated on them may have led to the delivery of volatile and organic species to render other worlds habitable.

Strategic Research for Q10.2

• Determine the extent of present and former habitable environments across the solar system by making measurements that determine the past and present existence of liquid water, the organic content, and the availability of nutrients and metabolic energy sources on terrestrial planets and ocean worlds.
• Determine the distribution of past and present subsurface oceans—fully liquid and muddy—and their historical evolution through detailed geological/geophysical investigations and modeling efforts coupled with a search for oceans by remote sensing.
• Determine the evolution of the climate of Mars and Venus and the timing of changes by measurement of atmospheric gases, chemistry, and isotopes in the atmosphere and rocks, and climate modeling.

Q10.3 WATER AVAILABILITY: WHAT CONTROLS THE AMOUNT OF AVAILABLE WATER ON A BODY OVER TIME?

Water can be mineral-bound in rocky materials or occur as ice, liquid, or atmospheric water vapor. The initial inventory of water for a planetary body is provided during and shortly after planetary formation, and its later availability is then a balance between subsequent supply and loss processes.

The presence of liquid water depends on distance from the Sun, the mode(s) of heat and volatile transport on a planet, and composition; for example, salts and volatile compounds (e.g., ammonia) can act as “anti-freeze” to lower melting temperatures. For the terrestrial planets, water availability in the crust or at the surface depends on the history of delivery during accretion, outgassing from the deep interior, physical properties of the subsurface such as porosity and permeability, loss to space, and transport and phase changes within and on the planet driven by temperature and pressure. For outer solar system objects with large volatile volumes, liquid water can be available beneath and within the ice shell, and in some cases in direct contact with rock, depending on the subsurface temperature and pressure conditions.

Q10.3a What Sets Initial Limits on Planetary Liquid Water and Water Inventories?

The initial water inventory depends on volatiles accreted during planetary formation by accretion and impact bombardment and is influenced by the position of the body in the solar system. Water is supplied from accreting planetesimals, asteroids, and comets, incorporated into coalescing objects or delivered later as a surface veneer (Question 3, Chapter 6). In the outer solar system, the strong gradient of water abundance in the Galilean satellites with distance from Jupiter suggests that water also can be accreted directly from a protosatellite disk surrounding giant planets.

Attempts to quantify the source(s) of Earth’s water and the history of accretion from the isotopic composition of its water and that in asteroids and comets have so far been inconclusive, given the large variations observed within measured planetary bodies and the uncertainty of how water is accreted into their source materials (see Question 3.6). In addition to accretion from the planetary nebula and subsequent aqueous alteration,

a contribution has been suggested for highly hydrated dust grains resulting from the interaction of the solar wind with planetesimals (Daly et al. 2021). Dynamical models of primordial water supply demonstrate the importance of planet formation location, planet size, and random variations in impacts by volatile-rich objects.

In the terrestrial planets and large moons, water can be released to the surface catastrophically during accretion and core formation. Liberated water can oxidize iron, releasing hydrogen; the hydrogen can then escape to space in hydrodynamic loss that can drag other gases with it (Zahnle et al. 2019). The difficulty in establishing robustly how much water has reacted with planetary materials, and thus how much water or hydrogen has been lost to space, challenges efforts to determine a priori how much water a planetary body was born with (Q6.1a).

Q10.3b What Are the Long-Term Endogenic and Exogenic Controls on the Presence of Liquid Water on Terrestrial Planets?

Aside from Earth, we have the most information to disentangle processes controlling availability of water on Mars (Figure 13-3). Major advances were made this decade thanks to MAVEN observations, integrated with compositional data from meteorite studies as well as orbiting and surface missions over two decades of observations that “follow the water,” but some questions remain at Mars, as well as for Venus.

The availability of liquid water at the surface of the terrestrial planets is determined in part by surface temperature, which itself is a function of distance from the Sun and atmospheric greenhouse warming. Although only small amounts of water are present in the clouds today, early Venus may have had a much more clement climate before a greenhouse effect drove liquid water into the atmosphere then away to space (Q6.2b). We know relatively little about early Venus because the mineralogy and chemistry of the rocks that record it are not known. On Mars, whereas the presence of abundant liquid water at the surface early in its history is undisputed (Q10.2), the mechanism for providing greenhouse warming to allow liquid water, either globally or locally, remains poorly understood (e.g., Wordsworth et al. 2017). Uncertainties include the amount of water present as ice or liquid within the martian crust both through time and today (Q.10.2), whether there was an early ocean, and the timing of water loss and climate change.

Volcanism supplies water and other gases to terrestrial planet surfaces and atmospheres through much of their history. The release of volatiles is directly correlated with the amount of volcanism, but the amount of water supplied depends critically on the amount of water incorporated during accretion, the redox state of the interior, and styles of volcanism, particularly temperature and degree of partial melting. Early voluminous volcanism on Mars and most other solar system bodies waned over time, resulting in progressively less water released, although net quantities are not well understood (Figure 13-3). Venus, too, experienced volcanism that drove outgassing, but activity prior to about 1 Ga is no longer preserved, making the history of volatile release there a key unknown.

In addition to endogenous supply from volcanism, impacts might have contributed to transiently thicker atmospheres that supported liquid water. Impacts producing craters hundreds of kilometers across and larger occurred for up to a billion years after planetary formation. Their role in water loss or gain and their environmental effects on atmospheric composition and climate remain key questions. These impactors delivered water and other volatiles to planetary surfaces, and impact heating and ejection may have liberated water ice or CO₂ volatiles. Yet, impact shocks may also have removed substantial atmospheric water and other volatiles. On the whole, the net role of large impacts in water loss or gain and their environmental effects on atmospheric composition and climate remains an open question.

Liquid water can be removed from a planetary surface by freezing, infiltrating the crust to become groundwater, or hydrating rocks (i.e., chemical weathering). On Mars, the volume of water ice beneath the surface and at the poles is equivalent to a ~40 m global-equivalent layer. Such ice could melt via geothermal-, volcanic-, or climate-driven heating. A more permanent means of removal from a hydrologic system is chemical weathering, especially if the planet lacks a means of crustal recycling such as plate tectonics. On Mars, hydrated minerals comprise a global equivalent layer of water 150 to 1,500 m thick (e.g., Scheller et al. 2021). On Venus, we do not yet know whether the crust includes hydrous phases (or their alteration products) that sequestered water or other volatiles.

Atmospheric escape processes deplete water from planetary atmospheres by a variety of thermal, photochemical, and charged particle interactions. Extreme ultraviolet solar radiation plays a role by the photodissociation of H₂O into component ions; the resulting hydrogen escapes to space as H or H₂. Ions also can be picked up by the solar wind and carried to space directly by sputtering, or the solar wind can accelerate molecules in the upper atmosphere and carry other atoms or molecules to space. Many of these processes have been observed directly at Mars and Venus and are inferred to have operated there over extended geological time, albeit at largely unconstrained rates. Calculated loss rates at Mars suggest a ~20 m global equivalent layer of water lost over time and possibly as much as ten times greater than that (Jakosky et al. 2018). The extent to which these processes affected the volatile inventory at Venus remains unclear. For Venus, the measured atmospheric deuterium–hydrogen ratio—more than 100 times that of Earth—is indicative of the loss to space of potentially substantial volumes of water (Donahue et al. 1982). Whether this elevated ratio represents a dynamic balance between supply of water by impacts with escape to space, the loss of steam from a primordial atmosphere, or a long-sustained vanished ocean driven away by a late-onset runaway greenhouse effect, remains a key open question.

Q10.3c What Are the Long-Term Endogenic and Exogenic Controls on the Presence of Liquid Water on Icy Bodies?

There is a strong coupling between the thermal history of icy bodies and the occurrence and persistence of subsurface liquid water. Oceans occur beneath icy shells because of heat supplied from planetary body cores, ultimately derived from accretional, radiogenic, and tidal heating, and potentially facilitated by the presence of

antifreeze in the oceans themselves. Ammonia and salts can facilitate melting by freezing point lowering, but an interior heat source is still required. Tidal dissipation accounts for subsurface liquid water oceans on some of the icy bodies, for example, via the Laplace resonance maintaining forced eccentricity among three of the Galilean satellites. A key question for Enceladus is the duration of its ocean, where interior cooling should outpace energy supply by tidal dissipation, unless the core is porous (Choblet et al. 2017). In contrast, Callisto, for example, experiences essentially no tidal energy dissipation, but it has a sufficiently high interior rock mass to sustain an ocean by radiogenic decay. Triton may have experienced considerable internal heating from tides during the early evolution of its orbit around Neptune. Most of the numerous known and suspected subsurface water oceans in the outer solar system are almost certainly in direct contact with rock, for example, at Enceladus and Europa, and thus could facilitate hydration of that rock. How far this liquid percolates into underlying rock depends on poorly known factors such as the porosity, permeability, and fracture density of the seafloor (Q10.6b,c).

For liquid water to reach the surface requires some means to overcome the higher density of liquid water relative to water ice. Cryovolcanism has also been suggested on Titan (e.g., Lopes et al. 2013) as an interpretation of surface features observed by Cassini, but on Titan—and Ganymede—where the subsurface oceans are trapped between two ice layers, the connection to either the silicate interior or the surface is not yet known. The only confirmed ocean–surface link is at Enceladus, where the composition of plumes emanating from south pole fractures indicates that they are sourced from a salty, subsurface ocean in contact with a rocky interior. By contrast, it is not clear that plumes on Triton are in contact with an ocean as those plumes could be modulated by seasonal solar irradiance and a solid-state greenhouse effect (Hansen et al. 2021). Hubble Space Telescope observations and a reanalysis of Galileo data have shown evidence, still debated, for plume activity on Europa (e.g., Paganini et al. 2020). A major question for Ceres is whether it presently has liquid water or brines beneath its surface, or if recently emplaced (~10 Ma) salts discovered by the Dawn spacecraft are mobilized remnants of an ancient ocean (Castillo-Rogez 2020). Liquid water may have been transiently present on many icy bodies following the heating associated with impact events, lasting for up to 10,000 years.

Strategic Research for Q10.3

• Establish whether liquid water is present on Mars today in the subsurface by geochemical measurements of ices and recent hydrous minerals and geophysical measurements to probe the upper crust.
• Determine the distribution, history, and processes driving the availability of ice and liquid water on Mars over time, combining mapping stratigraphy and mineralogy, measurements of chemical, mineralogic, and isotopic measurements in situ and from returned samples, sounding of the subsurface, models for geomorphic features and climate processes, and constraints on chronology from in situ radiometric dating and measurements on returned samples.
• Determine the availability through time of liquid water on Venus using measurements of present-day escape rates, isotopes and mineral phases in the crust, as well as atmospheric models integrating loss to space with interior and surface evolution.
• Identify the amounts and locations of any past or present liquid water beneath or emplaced on the surfaces of Enceladus, Europa, Titan, Ceres, and candidate ocean worlds within the satellite systems of Neptune and Uranus using radar, gravity, topography, magnetic field (induction), and surface spectral measurements, combined with models of tidal deformation and the formation and evolution of surface features.
• Determine the amount and origin of water ice on the Moon and Mercury by sampling ice, determining its spatial distribution, measuring H and O isotopes, and determining the nature and abundance of contaminants within the ice as a means of understanding sources of water in the inner solar system.

Q10.4 ORGANIC SYNTHESIS AND CYCLING: WHERE AND HOW ARE ORGANIC BUILDING BLOCKS OF LIFE SYNTHESIZED IN THE SOLAR SYSTEM?

Organic molecules are an essential component of planetary habitability, as they enable a variety of chemical structures and specific reactivity that is foundational to the biochemistry of life as we know it. Abiotic processes can produce diverse, large organic molecules, which are found in many places beyond Earth including Mercury,

Mars, chondrite parent bodies, Ceres, Enceladus, Titan, Pluto, comets, the atmospheres of the giant planets, and possibly Venus (Table 13-1). Of particular relevance to the study of habitable solar system environments are the recent discoveries of organic matter in water-deposited sediments on Mars, organic chemistry and prebiotic molecules on Titan, and organic matter in plume particles from the interior ocean of Enceladus.

To understand how abiotic organic synthesis occurs, it is important to have an inventory of organic products, reaction rates, and isotopic tracers, and their dependence on conditions in the present or past solar system. Such knowledge provides insights into prebiotic organic inventories that might factor into the origin of life, as well as which organic molecules can contribute to habitability by serving as carbon and energy sources for life. Moreover, a comprehensive assessment of abiotic organic processes will lead to an understanding of how signatures of biotic and abiotic products can be distinguished in the search for life on other worlds (Question 11).

Q10.4a Where Were Organic Compounds in the Early Solar System, How Did They Form, and Did They Contribute to Prebiotic Chemistry?

Organic molecules have been available since the birth of the solar system, formed abiotically via chemical processes. The high abundances of organic materials in carbonaceous chondrites, and especially in comets, show that the outer region of the early solar system was a site of organic molecule formation. These organics formed mainly via ion–molecule reactions at low temperatures in the outer protoplanetary disk and/or in presolar interstellar environments and include insoluble organic matter that is found in chondrites, as well as a diversity of smaller organic building blocks bearing nitrogen and oxygen atoms. These organics were transported in the disk (e.g., by turbulent mixing and gravitational scattering) and accreted by larger objects, including those in the inner solar system. Where exposed to warm liquid water—for example, in carbonaceous chondrite parent body asteroids, planet surfaces, and icy world interior oceans and early atmospheres—accreted organics underwent reactions that produced more complex compounds, including key building blocks of life such as amino acids, nucleobases, and sugars.

Q10.4b What Processes Have Enabled Abiotic In Situ Organic Synthesis and Cycling on Planetary Bodies?

Organic synthesis requires carbon and energy. In planetary surface environments, simple organic molecules like methane (CH₄) or formaldehyde (CH₂O) initially seed organic chemistry driven by ultraviolet light from the Sun, charged particles, and cosmic rays. In subsurface environments, the usual carbon source is CO₂; the key energetic driver is chemical disequilibrium with coexisting molecular hydrogen, and transition metals can catalyze synthesis.

Exploration of the solar system, as well as laboratory experiments and chemical kinetic modeling, have revealed how abiotic organic synthesis occurs in methane-bearing atmospheres. We have learned the most about this process on Titan, complemented by additional data from studies of Pluto, Triton, and the atmospheres of the giant planets. Abiotic organic synthesis in these environments is initiated by interaction of methane molecules with high-energy particles that enter the upper atmosphere, leading to the production of an array of radicals and electrically charged chemical species that combine in myriad ways to produce a diversity of organic molecules from simple gases to sooty aerosols (Figure 13-4). When nitrogen and oxygen sources are available, even more complex compounds can be formed by the incorporation of these elements into organic structures, leading to the formation of prebiotic molecules.

Direct sounding of Titan’s upper atmosphere by Cassini allowed for the detection of numerous species (e.g., Coustenis 2021). Higher molecular weight species (up to several 1,000 Da) are present in the ionosphere, including hydrogenated amorphous carbons and polycyclic aromatic hydrocarbons (PAHs), which can become long enough to form fullerenes. Methane exists in different phases on Titan, supporting a cycle similar to that of water on Earth and likely being replenished from the interior over long periods of time. The cycling of volatile chemicals starts with the dissociation of N₂ and CH₄ through electron, photon impacts, and cosmic rays higher in Titan’s atmosphere, leading to the formation of acetylene (C₂H₂) and hydrogen cyanide (HCN). Several other hydrocarbon compounds (e.g., C₂H₆, C₄H₂, C₂H₄, C₆H₆) and nitriles (e.g., HCN, HC₃N, C₂N₂,) were detected in Titan’s stratosphere from space or from Earth. Once formed, these molecules diffuse downward, forming haze and higher hydrocarbons and

nitriles, eventually depositing on the surface (see Figure 13-4). The extent to which they find their way to Titan’s internal ocean remains a matter of debate.

Analogous photochemical processes produce organic compounds on Triton and Pluto. Both worlds have N₂–CH₄ atmospheres, but are colder than Titan, so methane gas is less abundant relative to nitrogen. Pluto’s atmosphere contains C₂H₆, C₂H₄, C₂H₂, and HCN, and its surface also hosts organic compounds (e.g., Cruikshank et al. 2019). A key consequence of the different relative abundances of atmospheric N₂ and CH₄ is that the types of photochemically synthesized organic compounds on Triton and Pluto are predicted to be richer in prebiotically relevant nitriles than at Titan (Wong et al. 2015). The relative abundance of nitriles can constrain past levels of atmospheric methane.

Water–rock reactions are also key sources of solar system organics (Figure 13-5). Abiotic, seafloor H₂ and CO₂ reactions in hydrothermal systems on Earth produce organics, and there is potential for similar processes in the subsurface of icy ocean worlds, particularly Enceladus and Europa (Q10.5). Ice particles from the plumes of Enceladus contain silica as well as Na, Cl, and CO₃, derived from ocean waters reacting with rock in the core (Hsu et al. 2015). Cassini sounded Enceladus’s plumes and detected water vapor, molecular nitrogen and hydrogen, ammonia, carbon dioxide, and traces of several organic components: methane, propane, acetylene, benzene, and formaldehyde, similar to what is found in most comets. Importantly, complex macromolecular organics were identified in emitted ice grains (e.g., Postberg et al. 2018a) and amines are also indicated, providing key building blocks for life.

Abiotic organic synthesis at lower water–rock ratios, such as those of rock hydration, was important on Mars and some asteroids. Studies of carbonaceous chondrites suggest that such meteorites contain diverse organic compounds sourced from presolar origins that have undergone evolution in hydrothermal environments (e.g., Vinogradoff et al. 2018). Although difficult to characterize, the insoluble, macromolecular fraction of this material is composed of small PAHs (hydrocarbon rings) with abundant nitrogen, oxygen, and sulfur functions, cross-linked by short, highly branched carbon and carbon–oxygen chains. The soluble fraction comprises amino acids, chained amines (C–NH₂) and amides (C–N=O), chained and ringed hydrocarbons, alcohols and polyols (C–OH), aldehydes and ketones (C=O), chained carboxylic and hydroxy acids (O=C–OH), nitrogen-bearing rings and nucleobases, alkyl sulfonic (O=S=O), and phosphonic (−PO₃) acids, which are all of interest to astrobiology (Glavin et al. 2018) and may have contributed to the origin of life on Earth. The analysis of organics in carbonaceous chondrites gives insights on the prebiotic chemistry that occurred in a natural system before life. Telescopic surveys have confirmed infrared absorptions related to C–H in select large asteroids; more recently, the Dawn mission at Ceres found evidence for a high abundance of carbon in the regolith, interpreted as a mix of carbonates and amorphous carbon. Select locations on the surface are enriched in organic chains of atypical composition or abundance, thought to be generated from endogenous hydrothermal processes (De Sanctis et al. 2019).

On Mercury, large swaths of the crust contain graphitic carbon, inferred to be magmatic in origin. Mars also has indigenous organics (found in martian meteorites), and there are hints at a complex organics cycle involving (1) infall of carbonaceous meteorites and interplanetary dust particles; (2) primary igneous reduced carbon—that is, magmatic graphite; (3) primary hydrothermally formed organic carbon/nitrogen-containing species; (4) secondary hydrothermally generated macromolecular carbon and graphite; (5) impact-generated graphite; and atmospheric and aqueous processing of these (Steele et al. 2016). Select sedimentary rocks from a martian paleolake in Gale crater contain organic rings and chains bearing sulfur (Eigenbrode et al. 2018), suggesting transformation and

preservation in the presence of H₂S. So far, the martian organics can all be explained by abiotic processes, although our understanding of organic compounds on Mars is nascent, and strategies are being developed to discriminate between abiotic and biotic organics (Question 11).

Q10.4c How Have Organic Fluxes to Planetary Bodies Changed Over Time and Influenced Habitability?

Outer solar system planetesimals and protoplanets delivered both volatiles and evolved primordial organic material to growing worlds in the inner part of the system via impacts (Question 3). This delivery has generally decreased over geological time as leftover materials have been gradually swept up into planets or the Sun or scattered out into interstellar space. However, delivery continues, and there may have been punctuated events of increased supply during episodes of enhanced bombardment. The availability of these compounds and their concentrations in select environments could have played a role in the origin of life by providing building blocks for complex organic chemistry (Question 9).

Strategic Research for Q10.4

• Determine the complexity attained by organic chemistry in Titan’s atmosphere, its sources and sinks, and its role in producing a potentially habitable environment by entering the subsurface ocean, through in situ and remote spectral imaging and mass spectrometry investigations.
• Determine the chemical composition and structural characteristics of organic compounds in both Europa’s and Enceladus’s oceans to understand organic synthesis, using in situ techniques (chemical derivatization, gas chromatography, capillary electrophoresis, pyrolysis, and mass spectrometry) on plume or extruded materials, and detailed characterization by sample return and analyses in Earth laboratories.
• Understand the formation and alteration of organics on small bodies (asteroids, comets, meteorites, and dwarf planets such as Ceres) and their potential contributions to the origin of life on Earth or elsewhere by in situ observations of organics and isotopic composition, laboratory analog experiments, detailed investigation of samples returned to Earth and meteorites.
• Characterize organic molecules present on Mars for determination of type, distribution, and source of organic materials to understand organic synthesis there by in situ measurements and sample return for investigation in Earth laboratories.

Q10.5. WHAT IS THE AVAILABILITY OF NUTRIENTS AND OTHER INORGANIC INGREDIENTS TO SUPPORT LIFE?

Six main chemical elements make up life—carbon (C), hydrogen (H), nitrogen (N), oxygen (O), phosphorus (P) and sulfur (S), or “CHNOPS”—along with some “micro-nutrient” metals, for example, Fe and Mo (Question 9). H, C, O, and N are commonly available in terrestrial planet atmospheres, starting out as H₂O, CO₂, and N₂, and can be exchanged with nonatmospheric reservoirs. In contrast, sources of any P, S, and “micronutrient” metals required to sustain habitability are more commonly expected to be sourced from minerals present in rocky silicate materials.

Q10.5a What Are the Inventories, Forms, and Distribution of Life-Supporting Elements on Planetary Bodies?

On Venus, the available inventory of key life-supporting elements remains uncertain, in large part because of limited chemical data for the surface and atmosphere. From those few geochemical measurements made at the surface, Venus’s expansive lowland plains appear basaltic, and thus those flows and their physical weathering products are presumably potential sources of Mg, Fe, Ca, K, as well as P and S, which are typically present in basalts as minor constituents. Chemical alteration pathways on Venus are poorly understood, although they are likely dominated by sulfatization and oxidation (Zolotov 2019). Carbon and oxygen abound because the atmosphere is 96.5 percent CO₂; nitrogen and sulfur are also present (the latter as SO₂); H is absent but for trace amounts of

H₂O and H₂S. Whether there is phosphorus in the Venus atmosphere remains unclear: P sourced from volcanic eruptions could exist as P₄O₆ in the lower atmosphere, before being lofted to middle- and upper atmospheric levels and converted to PH₃; this latter phenomenon was recently tentatively observed. There is virtually no information regarding micronutrients in the Venus environment at present.

All of the essential life-supporting elements are present on Mars and have also been detected in forms accessible to life. Oxidized carbon is broadly available; 95 percent of the atmosphere is carbon dioxide, and localized carbonates have been detected by orbital and landed missions. Reduced carbon is also available in a variety of organic compounds detected in drilled samples by Curiosity in Gale crater, in Mars meteorites (Q10.4b), and perhaps locally in methane (Q10.2a). Nitrogen, long known to be present as N₂ in the martian atmosphere, is also present as fixed nitrate in surficial samples analyzed by Curiosity (Stern et al. 2015). H₂O is present in solid and gaseous forms today with abundant evidence for its presence as a liquid in the past. Phosphorus is enriched on Mars’s surface relative to Earth’s by about an order of magnitude, as measured in the dust and by every landed mission. Sulfates are present in the ubiquitous dust, commonly found in fracture fills, present in sand deposits near the north pole and even form thick surface deposits; sulfides and metals in other minerals are found in Mars’s igneous rocks.

Among the key major elements required for life (CHNOPS), all but N and P have already been identified on Europa and all but P in the plume of Enceladus (Hand et al. 2009; Postberg et al. 2018b). Given the abundance of P and metals in chondritic materials (assumed to be representative of the rocky interiors of these ocean worlds), however, neither is likely to be limiting to life on either Enceladus or Europa (e.g., Cable et al. 2020). On Europa, surface S is pervasive as magnesium sulfate salts, at least in part owing to an exogenous flux from Io’s volcanoes. On Enceladus, H₂S has tentatively been measured by Cassini in the plume gas (Postberg et al. 2018b).

Q10.5b What Processes Govern the Inventories, Forms, and Distribution of Life-Supporting Elements on Planetary Bodies?

A key factor that affects the availability of life-supporting elements on a world is their initial inventory, determined by the nature of accreted materials and how they were delivered. Because H, C, N, O, and S tend to form species that are solid only at low temperature (e.g., CO₂, NH₃, CO, CH₄, N₂, H₂S), these compounds were depleted in inner solar system solids and enriched in outer solar system materials (as observed in comets today). However, large-scale mixing in the protoplanetary disk helped to blur some of these differences (Questions 3 and 1, Chapter 4). Cometary species such as CO₂, NH₃, and organic matter can provide large inputs of carbon and nitrogen as well as CO, CH₄, N₂, and H₂S to planetary bodies.

The early evolution of a planetary body also factors into the contemporary availability of life-supporting elements. For example, the formation of a metal core sequesters iron, and CHNOPS elements can partition with the iron if conditions are sufficiently reducing. Iron in the silicate portion of a body is predominately in the divalent state, a key source of reducing power. Both early outgassing during a magma ocean phase and later volcanic outgassing can provide volatiles to the surface-atmosphere system. The nature and abundances of these outgassed volatiles depend on temperature, pressure, and compositional characteristics (e.g., oxygen fugacity) of the source region. Once at the surface or in the atmosphere, the volatile composition can be modified by photochemical reactions, escape to space, reactions with surface materials, and the activity of life.

Radiation-induced breakdown fundamentally alters the chemistry, and particularly the redox state, of C, N, S, P, and metals. The formation of highly oxidized species can contribute to the generation of potent sources of chemical energy for life (Q10.6). On the inner planets, the dominant source of radiation is ultraviolet light (Q10.3b). When water molecules are broken apart, H₂ and O₂ or hydrogen peroxide are produced. Mars provides an example where the escape of H₂ to space led to the net oxidation of the near-surface environment and production of other oxidized species, including, for example, sulfates, ferric iron, that are derived from the oxidation of volcanic or crustal reduced species (e.g., sulfides and ferrous iron). In the outer solar system, water radiolysis is mainly driven by charged particles in planetary magnetospheres. At Europa, the population of charged particles consists of electrons with energies of a few electron volts (eV) to MeV, and heavy ions (mostly S⁺ and O⁺) with energies between a few keV and 1 MeV. Direct radiolysis occurs to a depth of tens of centimeters, but its products are mixed to several meters by impact gardening. Further recycling processes like convection and subsumption

of ice may transport oxidants to Europa’s ocean (e.g., Hand et al. 2020). At Enceladus, radiolytic processing of surface ice is diminished by approximately an order of magnitude relative to Europa, and thus there may not be a comparably efficient mechanism for generating new oxidants at the surface.

Water–rock reactions, such as dissolution and mineral formation, exert a major influence on the abundances of CHNOPS elements in liquid water. A source for these elements is dissolution of primary rocks and minerals, in which they may only be present as minor constituents. Balancing supplies from water–rock reactions, minerals that sequester these elements include carbonates (for C), ammonium salts and clay minerals (for N), phosphates (for P), and sulfates or sulfides (for S). Once formed, the solubilities driving the bioavailability of elements in these minerals depend on the composition of the coexisting water, and particularly on its pH. H₂ in fluids can be generated by water radiolysis (Q10.6) or from oxidation of Fe and S during alteration (Q10.6) and react with carbon dioxide or carbonate in waters and rocks to produce organic molecules (Sherwood Lollar et al. 2021) (see Figure 13-5).

Critically, for water–rock interactions to occur in the first place, these two environments need to be in direct physical contact. Worlds that likely have high-pressure ices at the base of their oceans, such as Titan or Ganymede, may not host extensive interactions (except, perhaps, where localized heating is sufficient to melt overlying ice). The situation is more promising on Europa, where a silicate seafloor is almost certainly in contact with a subsurface ocean, and even more so within Enceladus—where ice grains containing sodium or silicon, diagnostic of water-rock interactions at depth, are ejected in the south polar jets (Figure 13-6).

Q10.5c How and Why Have the Inventories, Forms, and Distribution of Life-Supporting Elements Changed Through Time?

The availability of CHNOPS on a planet varies with time by mechanisms similar to those controlling the availability of water (Q10.3). Impacts and outgassing by volcanoes supply new gases with CHNOPS to the surface and supply new rocks for weathering. Water–rock reactions during weathering can liberate CHNOPS from rocks, for example, via mineral and glass dissolution, or lock them away in an unusable form, for example, by crustal hydration and formation of carbonate, sulfate, or nitrate minerals. On Earth, the bulk of the planet’s CO₂ inventory is stored as carbonate minerals in the crust; CO₂ stored there is unavailable for use by organisms until the carbonates dissolve in water or are subducted into the mantle and released via volcanism. It is not clear whether such crustal recycling mechanisms have existed on Venus or Mars. Escape to space can remove considerable volumes of atmosphere (Q10.3) with a variety of potential consequences: directly decreasing the abundance and availability of gases that might otherwise be bioavailable; changing climate, leading to phase changes that make material more or less bioavailable; and changing redox conditions. For example, the time-integrated impact of loss of H₂ to space from radiolysis of water molecules leads to a buildup of O₂, which can be an important source of energy for life in chemical reactions (Q10.6). Last, seasonal or diurnal cycles can affect the distribution and availability of life-supporting constituents, as on Titan, where the atmospheric composition varies with time—for example, the poles become enhanced during hemispheric winter, as on Earth, leading to more complex molecules (Coustenis 2021)—and on Mars with seasonal cycles in atmospheric H₂O, O₂, CO₂, and methane, as observed by Curiosity. Such cycling can generate disequilibrium, for example, subsurface buildup of reduced gases from water–rock reactions and periodic release into a more oxic atmosphere.

Strategic Research for Q10.5

• Determine if life-supporting chemical species, including reduced carbon- and phosphorus-bearing molecules, are present in the atmosphere of Venus, via in situ atmospheric measurements.
• Determine the diurnal and seasonal variability of Titan’s atmospheric and surface chemical composition by long-term monitoring from space and in situ measurements with support from ground and Earth-bound observatories and modeling.
• Constrain CHNOPS speciation in liquid water environments on Europa, Enceladus, and other ocean worlds where materials sourced from an ocean can be accessed, using measurements of tracer species in plume gases and particulates, volatiles in the ambient exosphere, and surface mineralogy/composition associated with materials sourced from the ocean.
• Determine the inventory and bioavailability of CHNOPS, particularly reduced carbon and fixed nitrogen, for candidate current and ancient habitable environments on Mars by measurements of mineralogy, chemistry, and isotopes from high-resolution remote sensing, surface missions, and returned samples as well as modeling exchanges.
• Quantify the impact of the seasonal cycles in the martian atmosphere on the formation and long-term evolution of CHNOPS species, using telescopic, spacecraft, and in situ observations of the behavior of aerosols, trace gas abundances, and isotopes, as well as observations and models of meteorological behavior and surface fluxes of gases.

Q10.6 WHAT CONTROLS THE ENERGY AVAILABLE FOR LIFE?

Habitability requires that the environment supply energy to the biological system to meet organism demand. Life on Earth utilizes a variety of energetically favorable reactions for maintenance and growth, and there are even examples of community metabolisms where organisms employing families of redox reaction couples exist symbiotically (Question 9). The nature of energy sources available for life, whether there is adequate power (i.e., energy per unit time), and the longevity of energetically favorable conditions (e.g., long-term maintenance of fluxes that sustain disequilibria) are key considerations for dynamic habitability.

Q10.6a What Are the Available Energy Sources for Life?

All life as we know it exploits energy via thermodynamic disequilibria. The presence or absence of disequilibrium conditions was initially considered as a basis for a life-detection strategy. We now recognize that the relationship between life, energy, and thermodynamic disequilibrium is more nuanced. Although thermodynamic disequilibrium can cause life to emerge, promote the growth of living organisms, sustain extant life, and be generated as a byproduct of life, such disequilibria also exist independently of life (NASEM 2019a).

Three primary energy sources power reactions that can support life on a planetary body:

• Solar radiation: Light is a dominant source of energy for life on Earth via photosynthesis. In oxygenic photosynthesis, photoreactive molecules are excited by absorbed light at wavelengths characteristic of a pigment (e.g., chlorophyll), triggering biochemical reactions that split water molecules, transform CO₂ to carbohydrate, and release O₂ as a by-product. In anoxygenic photosynthesis, different electron donors (e.g., hydrogen sulfide) are used instead of water for reactions with carbon dioxide. On Earth, oxygenic photosynthesis evolved relatively late, but anoxygenic photosynthesis likely sustained the earliest photosynthetic life forms, including early Archean shallow-water stromatolites (Q9.3). The discovery of diverse pigments exploiting wavelengths from the visible to near-infrared and many electron donors (Kiang et al. 2007) indicates that photosynthesis could be a viable energy source in a wide range of solar-illuminated planetary surface environments, although a challenge elsewhere in the solar system is meeting the other conditions for surface habitability.
• Chemical energy: The biomass of chemotrophic organisms on Earth is comparable to the photosynthetic biosphere, sustained by energetically favorable chemical reactions that support a variety of aerobic and anaerobic metabolisms: sulfide and iron oxidation; methanotrophy; methanogenesis; acetogenesis; nitrate, sulfate, and iron reduction; and fermentation (Onstott et al. 2019). These metabolisms are widespread in hydrothermal and groundwater systems, at volcanic vents, and in lakes and oceans at surface-atmosphere and water-bottom interfaces. Chemical energy sources are viable on other worlds; simple, widely available species participate in energy-supplying redox reactions, including H₂O, H₂, O₂, CO₂, CH₂O, CH₄, organic C, SO₄²⁻, HS⁻, NO₂⁻, NH₄⁺, and Fe^2+,3+. Environments with these redox reactions existed and may yet exist on Mars, water-altered asteroids like Ceres, and ocean worlds including Europa, Enceladus, Ganymede, and Titan.
• Radiolysis: Exposure to radiation breaks chemical bonds, creating new species that supply chemical energy for the redox reactions above. In endogenic radiolysis, radiation from radioactive uranium, thorium, and potassium isotopes in rock breaks apart water molecules in rock-permeating fluids to generate reductants (H₂) and oxidants (H₂O₂). Subsequent redox reactions involving these products sustain life kilometers underground, largely disconnected from other water reservoirs. By analogy with Earth, calculations of the energetic fluxes provided by radiolysis in the subsurface on Mars (Tarnas et al. 2021) and ocean worlds (Bouquet et al. 2017; Altair et al. 2018) suggest that those fluxes may be sufficient to sustain organisms. On airless bodies, additional exogenic radiolysis arises from solar irradiation or magnetospheric plasmas. For example, particles incident from Jupiter and from solar irradiation break water molecules at the surface of Europa’s ice shell. H₂ is generated and typically lost to space, enriching oxidizing species such as O₂ and H₂O₂, and setting up a redox disequilibrium with materials in the interior. An outstanding question is whether these materials cycle back into Europa’s subsurface ocean.

Q10.6b What Geophysical and Geochemical Processes Determine Power Available for Life?

Geological processes on a planetary body drive energy availability. Some sources of energy are crucial in sustaining elements of habitability but do not themselves directly act to sustain life. For example, tidal energy sustains liquid water oceans in the subsurface of icy satellites, and heat from geothermal and hydrothermal energy can maintain liquid water in the subsurface of Mars and perhaps even Ceres today. For the sources of energy critical for life directly, the major question is whether the environmental energy flux is sufficient to support life. Voltage (i.e., energy per reaction quantum) and power (i.e., energy per unit time) are the key parameters to not

only sustain life in a dormant state against environmental stressors but to support metabolism, mobility, growth, and reproduction (Hoehler 2007). The energy for life (J/mol) available in planetary environments can be calculated, given sufficient compositional data, from fluid chemistry (at the present) or mineral assemblages (in the past). The power for life, or timescale of the energy supply (J/s or W), depends on the fluxes and the energetic reactions available (Q10.6a). There is a minimum maintenance power to keep the cells viable, probably between 1–1,000 × 10⁻²¹ W/cell (LaRowe and Amend 2015).

For photosynthetic metabolisms, power is based on the luminosity of the Sun at the orbital radius of the planetary body, accounting for attenuation by an atmosphere and for diurnal, seasonal, and longer-term (e.g., precessional) cycles. In the past, Mars’s surface may have had punctuated periods of habitability (Q10.2a) to support photosynthetic life; key questions include how long these surface habitats persisted and whether life took advantage of these settings (Q10.7). To what extent evolutionarily advantageous energy storage or periods of dormancy could take advantage of light energy on bodies spinning slower than Earth’s 24-hr cycle (e.g., the 243-day cycle of Venus) remains unclear. More importantly, surface environments where photosynthetic energy for life is readily available face other habitability challenges. Liquid water availability is the challenge on Mars, availability of organic matter and water and very low pH is the challenge in Venus’s upper atmosphere, and availability of solvents for non-water-based life is the challenge for Titan.

Investigating whether energy is the limiting factor for life in potentially habitable subsurface environments found throughout the solar system is a priority (NASEM 2019a). At Europa and Enceladus, a key question is whether disequilibrium between water in their oceans and their rocky cores persists despite billions of years of chemical reactions. Alternatively, radiolysis (Q10.6a) may supply necessary redox couples to power life (Ray et al. 2021). At Enceladus, evidence for thermodynamic disequilibrium consistent with such processes has already been reported (Hsu et al. 2015; Waite et al. 2017). There, CO₂, H₂, and CH₄ co-occur in abundances where a strong energetic drive exists for CO₂ and H₂ to combine to form more CH₄ and water. Enceladus’s energy supply—which is likely the limiting factor, rather than liquid water or nutrients—could sustain a biomass density comparable to that of subglacial lakes on Earth (Cable et al. 2020). In the inner solar system, fluid circulation in rock was prevalent in the past on Mars and Ceres, and may still occur on both, as well as on other large hydrous asteroids. On Mars, silicate hydration, iron oxidation, sulfate formation, and several other chemical processes operated in the past to supply energy to subsurface life (e.g., Onstott et al. 2019). Radiolytic energy appears to be sufficient to power life (Q10.6a), and there may even be modern hydrogen or methane fluxes locally (Moores et al. 2019).

In situ measurements or sample return are required to gain data on reactants and products for quantitative constraints. Then, modeling of anticipated water-rock reaction conditions and their timescales offers potential for exploring the past and present habitability of other planetary bodies, albeit with caveats. First, theoretical modeling is most robust for situations in which equilibrium is anticipated to have been obtained, whereas, in a dynamic setting, reaction kinetics might dominate; this necessitates complementary study of Earth analogs and laboratory experiments (Question 9). Second, the maximum chemical energy available from any set of reaction conditions does not necessarily imply that life is able to exploit it, as other habitability conditions have also to be met. Third, the supply of energy is highly heterogeneous on a planetary body, and, therefore, so would be the distribution of any biomass. Focused reaction zones are where redox chemistry occurs, and the flux of reactants and products is controlled by permeability, fracturing, and hydrological or geophysical processes that refresh products and reactants. “Follow the interfaces” is a guiding principle in Mars exploration of the subsurface for life (Onstott et al. 2019). This approach is also applicable to ocean worlds where shell-ocean and ocean-seafloor interfaces are expected to be the zones with concentrated energy fluxes that could provide favorable habitats (see Figure 13-6).

Q10.6c What Processes Govern Energy Availability Over the Long Term?

Over geological timescales, the sustained release of chemical energy requires mechanisms for renewing reactants to prevent equilibration. In addition to radiolytic oxidant supply (Q10.6a, Q10.7b), these mechanisms include fracturing and erosion to expose fresh surface materials to chemical reaction; impacts, tectonic activity, and magmatism/volcanism to emplace or redistribute material out of equilibrium with its surroundings; and

hydrological parameters that control fluid circulation. In some cases, such processes might become transiently self-sustaining: for example, volumetric expansion associated with mineral hydration, as observed in serpentinization systems, can lead to fresh fracturing and, hence, further introduction of fluid. Presently, the rates of most such processes are unknown on most planetary bodies.

Strategic Research for Q10.6

• Geochemically characterize past and present environments with liquid water to determine whether there is/was energy to sustain metabolic processes of life by in situ or sample analysis of waters and preserved water-formed mineral and chemical species to determine concentrations of major ions, electron donors and acceptors, mineral products, and other relevant chemical species.
• Characterize the compositional and geological heterogeneity of potentially habitable worlds at progressively smaller scales (km- to cm-scale) to identify locales where chemical energy is or was more available by identification of mineralized fractures, reaction fronts, permeability boundaries, sites of ocean-ice exchange, and other interfaces via orbital remote sensing and in situ landscape- and microscopic-scale characterization.
• Determine the geophysical parameters that control past and present material fluxes in rocky subsurfaces, such as porosity, permeability, heat flux, volcanic flux, and tectonics by geophysical measurement, drilling/coring, change-detection experiments, seismic experiments, and modeling.
• Determine the kinetics of chemical reactions relevant to energy supply and material availability for life under conditions (past or present) on planetary bodies in the solar system by conducting laboratory experiments.

Q10.7 WHAT CONTROLS THE CONTINUITY OR SUSTAINABILITY OF HABITABILITY?

The continuity of habitable conditions is likely important for both starting and sustaining life on other worlds. If periods of habitability are too short to allow life to take root, a body or environment may never be capable of supporting life, even if it meets the definition of “habitable.” On the other hand, environments presently on the edge of habitability (e.g., Venus’s atmosphere) could sustain life that started under more clement conditions (e.g., on the Venus surface during an earlier habitable phase) or provide changing conditions conducive to the emergence of life, as has been suggested for episodically wet/dry surface environments on early Earth and Mars.

Q10.7a What Exogenous Factors Control the Continuity of Habitability?

Exogenous factors that affect planetary habitability include (but are not limited to) the properties of the host star(s), the orbital dynamics of the planet, and the impact bombardment flux (see Figure 13-2).

Worlds with slightly elliptical orbits such as Mars are exposed to a wider range of stellar fluxes over the course of a year than those with near-circular orbits such as Venus. Orbital resonances can provide heat and energy for ocean worlds such as Europa and Enceladus through tidal heating. A natural satellite that is large relative to its planet, such as Earth’s Moon, can stabilize obliquity (spin axis tilt) and climate; by contrast, Mars, with only relatively small satellites, experiences larger variations leading to ice being redistributed between polar and equatorial regions, and perhaps liquid water becoming transiently stable, over thousands to millions of years because of periodic variations of tens of degrees in axial tilt (Q10.2).

The relative contribution of impact bombardment to degrading or enhancing habitability on any given body remains to be fully understood. Extremely large impacts can strip planetary atmospheres or melt lithospheres, destroying prevailing habitable environments as well as any record of their presence. However, bombardment also delivered volatiles to the inner solar system, including Earth, that were critical to generating our hydrosphere and atmosphere (Q10.3). Impacts can destroy surface habitable environments (e.g., Earth’s K-T impact, which caused mass extinctions), but impacts can also simultaneously maintain or even create subsurface habitable environments, such as possible transient hydrothermal systems on Mars, Ceres, and Titan. The relative contribution of impact bombardment to degrading or enhancing habitability on any given body remains to be deciphered.

Q10.7b What Endogenous Factors Control the Continuity of Habitability?

The duration of habitability may be most strongly affected by the processes associated with heat transfer and loss over time, which are themselves a function of planetary size and starting composition (e.g., Ehlmann et al. 2016). For worlds that are sufficiently large and differentiated to have an intrinsically generated dynamo, a magnetic field could provide adequate protection of the surface and atmosphere from stellar activity and atmospheric stripping, or could alternately contribute to atmospheric loss (e.g., Gunell et al. 2018).

The tectonic regime of the terrestrial bodies and icy shells (Question 5, Chapter 8) can play a major role in establishing, sustaining, or destroying habitability. Surface mobility enables the exchange of nutrients, energy, and water between the interior and exterior. Numerous tectonic regimes have been proposed, including active-lid (a subset of which defines Earth’s present plate tectonics paradigm), heat pipe, sluggish-lid, episodic-lid, plutonic-squishy lid, and stagnant-lid modes, which differ in how heat is lost and how tectonic and volcanic activity is manifested (e.g., Lourenço et al. 2020; Byrne et al. 2021). This, in turn, determines how volatiles are cycled between the interior, surface, and atmosphere or sequestered, drawing down the atmosphere over time (Question 5; Q10.3). Bodies may transition between tectonic regimes over their lifetime (e.g., Weller and Lenardic 2018). Internal processes that redistribute mass can change a body’s orbital and rotational dynamics (e.g., the Tharsis Rise on Mars), as well as elemental availability and oxidation state, hence altering conditions in both surface and subsurface environments.

A tectonic regime conducive to long-term volcanic activity, as is likely the case for Venus, can replenish an atmosphere that might otherwise be lost to weathering or escape. Bursts of volcanic activity can have major effects on climates on short and long timescales (e.g., the Permian-Triassic mass extinction event on Earth). Volcanic activity likely helped Earth transition out of a past snowball state and may have precipitated Venus’s runaway greenhouse atmosphere and loss of surface habitability (Way and Del Genio 2020). On ocean worlds, material exchange between the surface, ice shell, and any underlying ocean is still poorly understood, as is the role of such cycling on habitability by promoting exchange of materials.

Strategic Research for Q10.7

• Assess surface–ocean exchange and the dynamics and long-term evolution of oceans and ice shells on ocean worlds, including Europa and Enceladus, by measuring the chemistry and speciation of plumes and surfaces with orbital remote sensing, in situ, or sample data, and by models that couple orbital and internal evolution and dynamics.
• Determine the nature, timing, and processes controlling the existence of past habitable environments on Mars by measurements in situ and in returned samples of stratigraphy, petrology, organic content, isotopes in rock, and geochronology in multiple environments covering multiple time periods.
• Identify the effects of large impacts on local and planetary habitability by investigating impact sites on potentially habitable bodies, analyzing chemical and isotopic signatures in rock/ice records before and after impacts, modeling the thermal and compositional effects of impacts on planetary atmospheres and surfaces, and improving knowledge of impact flux with time.
• Understand the diversity and controls on rates and styles of recycling of surface materials by remote sensing of evidence of these processes on relevant worlds, and by developing models for lithosphere dynamics and (cryo)volcanism that account for the thermal and orbital evolution of bodies, and the rheological and compositional evolution of their interiors.

SUPPORTIVE ACTIVITIES FOR QUESTION 10

• Improved radiative transfer modeling and photochemical modeling in planetary atmospheres for models of climate and organics production by fundamental laboratory measurements and computational work to obtain photochemistry reaction coefficients and gas absorption parameters.
• Improve the characterization of diverse ices and the mineral and organic products of water–rock interaction that are fingerprints of habitability by fundamental laboratory measurements to obtain and compile libraries of ice, mineral, and organic spectra and optical constants at ultraviolet to far-infrared wavelengths.

• Develop technologies for in situ measurement of light and radiogenic isotopes that trace the fluxes of key CHNOPS species and geological evolution in absolute time by miniaturization and maturation of instruments for isotope measurement on landed missions to terrestrial and ocean worlds.
• Increasingly high-fidelity analyses of organics and metal isotopes that determine abiogenic versus biogenic organic production as well as reservoirs and fluxes of key elements for biology by development of sample return technology for silicate, ice, and atmospheric samples and advanced facilities for returned sample analysis and curation.
• Determine the chemical and mineralogical outcomes of weathering, other aqueous alteration, and abiotic organic synthesis under non-Earth conditions by reaction path kinetic and thermodynamic modeling and with experiments.
• Characterize the limits of habitable conditions for solar system bodies by field work in terrestrial extreme environments, experimental studies, and modeling.

REFERENCES

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