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Question 6: Solid Body Atmospheres, Exospheres, Magnetospheres, and Climate Evolution

The nebular accretion processes outlined in Questions 1–3 (Chapters 4–6) gave rise to a diverse array of planets and moons that eventually evolved into the ones that we observe today.¹ In this chapter, the committee focuses on the atmospheres of these bodies, specifically those with solid surfaces, including the four terrestrial planets, the various moons, and the dwarf planets (e.g., Pluto). Gas and ice giant planets are discussed in Question 7 (Chapter 10). The atmospheres discussed here include dense atmospheres (e.g., Venus, Earth, Titan, and early Mars), atmospheres dominantly controlled by vapor-pressure equilibrium (current Mars, Triton, Pluto, and some other Kuiper belt objects, or KBOs), and collisionless exospheres (e.g., Mercury and the Moon), shown in Figure 9-1. Collisional atmospheres also have exospheres that separate them from the near vacuum of space. In the remainder of this chapter, the term “atmosphere” includes these exospheres. The study of planetary atmospheres is critical to understanding past and current habitability throughout the solar system, including the prebiotic processes that led to the emergence of life on early Earth. It also provides natural laboratories that we can use to better understand the processes governing Earth’s past and current climate.

How did the atmospheres of solid bodies form, and why did some of them end up with dense atmospheres while others did not? What processes contributed to atmospheric accumulation, and what processes led to atmospheric loss? How do climates evolve on solid bodies with atmospheres, and were any of them besides Earth capable of supporting life? This last question is revisited in more detail in Questions 9–11 (Chapters 12–14) of this report, but the foundations for that discussion are laid here.

Atmospheres also change on a variety of shorter timescales ranging from daily to seasonal to those encompassing orbital (Milankovitch) cycles. Some changes are closely linked to variations in solar forcing, while others (such as dust storms on Mars or changes in SO₂ abundance at cloud deck levels on Venus) are episodic and hard to predict given current understanding of these atmospheres. Long-term observations are needed to make progress toward understanding the processes responsible for such changes. Further changes (such as polar layered deposits on Mars or the migration of methane lakes between hemispheres on Titan) occur on much longer timescales, yet insight can still be gained by measurements of current conditions or through changes observed over several decades.

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

Many mysteries also have yet to be solved regarding the present atmospheres of solid planets and moons. What accounts for the ultraviolet absorption in Venus’s clouds and the reported occurrence of methane in Mars’s atmosphere? What is the detailed composition of the thick organic haze on Saturn’s moon Titan, and how does liquid methane cycle between Titan’s lower atmosphere and surface/near-subsurface? How do volatiles migrate on bodies with atmospheres dominated by condensation-sublimation flows, such as Triton and Pluto? How and to what extent do volatiles migrate within exospheres, such as on Mercury and the Moon? And what drives atmospheric superrotation on bodies like Venus and Titan, or dust storms on Mars and Titan? How do atmospheres interact with the space environment, and how is this interaction mediated by the presence of ionospheres and magnetospheres?

These are but a few of a fascinating array of questions that remain to be answered about the atmospheres of the planets and moons in today’s solar system. The committee highlights the major outstanding issues below.

Q6.1 HOW DO SOLID BODY ATMOSPHERES FORM AND WHAT WAS THEIR STATE DURING AND SHORTLY AFTER ACCRETION?

The formation of solid body atmospheres depends on an array of processes, including supply and removal of volatiles by impacts, atmospheric escape, and exchange with the interior. Many of these processes remain poorly understood. All solar system atmospheres have evolved significantly since their formation; thus, to investigate their earliest states we need to gather clues from a wide variety of sources.

In the inner solar system, the major differences in the atmospheres of Mercury, Venus, Earth, and Mars arise in large part from differences in the planets’ masses, distance from the Sun, and the presence or absence of a significant magnetosphere. Mars may also have formed as volatile-rich, but because it is small, it has lost much of its atmosphere to space over time. Mars’s small size may have also affected the chemical composition of its mantle, and hence its early atmosphere (Wade and Wood 2005; Deng et al. 2020), with important implications for climate (Q6.2).

Farther out in the solar system, the rich atmospheric diversity of the icy moons and dwarf planets poses fascinating challenges in comparative planetology. Titan stands out because of its thick nitrogen and methane atmosphere, in stark contrast to similarly sized objects such as Ganymede, which has only a tenuous oxygen exosphere. Understanding why Titan has retained a thick atmosphere while other outer solar system bodies did not remains a key challenge. Pluto and Triton, similarly sized bodies, have N₂ atmospheres in vapor pressure equilibrium with surface ices but the structure of their atmospheres is quite different. The oxygen exospheres of Europa and Ganymede are likely produced by sputtering of surface ice, but our understanding of this process has yet to be validated.

Progress in understanding atmospheric formation across the solar system requires a range of approaches, but from an observational standpoint, obtaining more precise isotopic data on volatiles and noble gases in solid body atmospheres, chondritic meteorites, and comets is particularly important. Noble gases—particularly He, Ar, and Xe—have both radiogenic and nonradiogenic isotopes that can be used to discriminate between early and late outgassing of a planet’s atmosphere (Avice and Marty 2020).

Q6.1a What Was the Role of Accretion and Meteoroid Impacts in Sculpting Early Atmospheres?

The earliest atmospheres of the inner solar system planets formed while the planets were still accreting, and the accompanying impacts both delivered and ejected volatiles. The volume and flux of impactors, their volatile content, and their degree of differentiation are all critical to determining what types of atmospheres first formed on the terrestrial planets. At the Moon, asteroids are believed to have delivered water during the lunar magma ocean period, suggesting that they also delivered water to early Earth (Barnes et al. 2016). Recent work suggests that transient, hydrogen-rich atmospheres caused by the thermochemistry of large impacts may have been common on planets like Earth and Mars (Haberle et al. 2019a; Zahnle et al. 2019). This process has important implications for early habitability (Q6.2) and prebiotic chemistry (Question 9; Question 11), but it remains poorly understood.

Q6.1b What Was the Role of Hydrodynamic Escape in Early Atmospheric Evolution?

Early atmospheres of planets and moons were influenced by many of the same thermal and nonthermal escape processes that occur today. (See Q6.5 for a discussion.) But early atmospheres were also likely influenced by an additional escape process, hydrodynamic escape, that is not observed on present solar system bodies, with the possible exception of Titan (Strobel 2009; Schaufelberger et al. 2012). Hydrodynamic atmospheric escape occurs when the upper regions of a planet’s atmosphere are dense enough that the pressure force remains important at all altitudes. The frequency of collisions also ensures that this is a form of thermal escape. In the extreme case, the density is high enough that the atmosphere can expand outward to space as a collisional fluid, so the process can be simulated using the standard equations of hydrodynamics, albeit formulated in such a way as to be able to handle transonic flow (Johnstone et al. 2018). The high levels of extreme ultraviolet radiation (XUV) produced by

the young Sun would have made this process particularly important to atmospheric evolution in the first few hundred million years after the planets formed (Lammer et al. 2018). This process likely drove loss of the primordial atmospheres of Earth, Venus, Mars, and Titan, setting the stage for later atmospheric and climate evolution (Q6.2).

When one studies this process in more detail, one finds that hydrodynamic escape grades continuously into a more familiar form of thermal escape, Jeans escape, that operates on multiple bodies today (e.g., Earth and Mars). In Jeans escape, the atmosphere is modeled as being collisionless above some level, termed the exobase. Ballistic particle trajectories are assumed above the exobase, and the atmosphere is assumed to be hydrostatic below it. But there is an intermediate regime in which an exobase occurs, yet the bulk atmosphere is not hydrostatic, and the pressure force is not negligible above it. In this regime, more complicated kinetic treatments of the Boltzmann equation are needed to accurately calculate escape rates (Johnson et al. 2013).

Hydrodynamic escape causes fractionation and hence can be investigated empirically by making precise measurement of isotopic variations in atmospheric noble gases, especially Xe (Zahnle et al. 2019). Interpreting such data requires a better understanding of solar extreme ultraviolet variation through time and of upper-atmospheric chemical and radiative processes. Additional modeling work considering multiple species and isotopes and comparing different numerical approximations are needed to understand thermal escape of early atmospheres.

Q6.1c What Was the Role of Magma Oceans in Early Atmospheric Evolution?

During formation, the surfaces of the terrestrial planets were likely molten owing to heat from accretion and blanketing by thick atmospheres, leading to formation of a local or global magma ocean (Elkins-Tanton 2012). Even when the magma ocean lifetime is short in geologic terms (tens to hundreds of million years), its influence on subsequent atmospheric evolution can be profound. For example, Earth’s mantle is thought to be more oxidized than that of Mars because of processes that occurred within its deeper, higher-pressure magma ocean (Wade and Wood 2005; see also Figure 9-2). Consequently, Mars’s early atmosphere may have been more H₂-rich than Earth’s

early atmosphere, and this may help explain why Mars’s early climate was relatively warm despite the faintness of the young Sun (Ramirez et al. 2014). Measurements of noble gases in planetary atmospheres and crustal samples are needed to constrain the earliest outgassing rates (particularly for Mars, Venus, and Titan), while laboratory experiments are needed to determine the chemistry and physics of interior exchange processes at temperatures and pressures relevant to magma ocean conditions.

Q.6.1d What Role Does the Space Environment Play in Forming and Liberating the Volatiles Contained Within Surface Bounded Exospheres Like That at the Moon and Mercury? What Role Do the Magnetospheres of Jupiter and Saturn Have on the Development of Exospheres at Ganymede, Europa, Dione, and Rhea?

On bodies with surface boundary exospheres, the solar wind and micrometeoroids are directly incident on the surface and are capable of altering surface material to create exospheric volatiles and to influence the cycling of the released volatile species. For example, the solar wind is directly incident at the Moon for three quarters of its orbit and implants protons in the top ~30 nm of grains. This proton implantation has been suggested to be the origin of a thermally modulated hydroxyl signature observed in the infrared (Li and Milliken 2017). Water has been observed to be released from the lunar surface during meteor streams and is believed to be manufactured from the solar wind implanted hydroxyls when the local surface near an impact is undergoing “flash” heating (Benna et al. 2019). It remains unclear if this transient water exosphere can migrate to the poles. Molecular hydrogen is also emitted from the lunar surface as part of the solar wind proton/hydrogen diffusion process (Hurley et al. 2016; Tucker et al. 2019). The solar wind-implanted hydrogen and carbon atoms may find each other via surface diffusion to form the observed methane emitted during the surface warming at dawn (Hodges 2016). We anticipate similar complex surface-exosphere interactions at other solid bodies that are directly exposed to the space environment, such as Mercury, or Phobos. We are still advancing our understanding of the complex surface-exosphere interactions for moons directly exposed to the energetic magnetoplasma environment of their parent planets such as Ganymede, Europa, Dione, Rhea. Both planetary and moon magnetic fields will also affect the processes (see also Q6.5, Q5.5a,c).

Q6.1e Was There an Early Collisional Atmosphere on Exposed Solid Bodies Like the Moon, Mercury, and Europa?

It has been hypothesized that the Moon had a collisional atmosphere shortly after the Moon-forming event during the lunar magma ocean period (Stern 1999). Needham and Kring (2017) further suggested that transient collisional atmospheres formed episodically in response to lunar mare volcanism. Some of the released mantle material may have been cold-trapped and sequestered within lunar polar crater deposits, especially the Cabeus region (Siegler et al. 2016). It remains unclear why the collisional atmospheres at the Moon dissipated. Similarly, it is unknown whether bodies that lack collisional atmospheres today, like Mercury or Europa, had significant collisional atmospheres early in their evolution and, if so, how those initial atmospheres dissipated over time.

Strategic Research for Q6.1

• Constrain the earliest stages of atmospheric evolution on Venus, Mars, and Titan by measuring noble gas abundances and isotopic fractionation to sufficient precision to quantify their minor isotopes.
• Derive the sources of exospheric volatiles by measuring the distribution, composition, and abundance of surface volatiles (including in permanently shadowed regions) on solid bodies including the Moon, Mercury, Ceres, and outer planet satellites such as Europa.
• Improve understanding of the initial states of planetary atmospheres by developing models of magma oceans coupled to processes such as delivery, loss of volatiles from impacts, chemistry, dynamics, and atmospheric escape owing to both thermal and nonthermal escape mechanisms.

• Develop physical and chemical constraints on early atmosphere–interior exchange by performing laboratory experiments on volatile partitioning in silicate melts, meteorite shock chemistry, and related phenomena.

Q6.2 WHAT PROCESSES GOVERN THE EVOLUTION OF PLANETARY ATMOSPHERES AND CLIMATES OVER GEOLOGIC TIMESCALES?

Planetary (and satellite) atmospheres probably formed during accretion, but they have continued to evolve over time as volatiles are added by volcanism or impacts or lost by escape to space. Volatiles can also be exchanged with a planet’s surface, especially on bodies like Earth or Titan on which liquids are present to facilitate erosion and weathering. Earth is also inhabited, so its atmospheric composition is affected by biogeochemical cycles and, more recently, by anthropogenic inputs.

Planetary climates are linked tightly to atmospheric density and composition. For the three terrestrial planets with substantial atmospheres, this is sometimes called the “Goldilocks problem”: Why is Venus too hot, Mars too cold, and Earth just right? Part of the answer clearly lies in their relative distances from the Sun. But a planet’s surface temperature also depends on the greenhouse effect of its atmosphere: clouds and hazes cool the surface by reflecting incoming sunlight, while various gases (e.g., CO₂ and CH₄) warm the surface by retarding the emission of thermal-infrared radiation. The density and composition of terrestrial planet atmospheres depend on a host of factors, including planetary size, volcanic activity, mantle redox state, and surface weathering processes. Exploring those factors can yield useful insights into why Earth remained habitable while its nearby neighbors did not and can help us understand whether Earth-like planets may exist elsewhere in the galaxy.

Outer solar system moons (Titan) and dwarf planets (Pluto) have their own volatile cycles that are controlled by a different set of processes. Titan, with its dense N₂-CH₄ atmosphere, has a climate analogous to Earth’s, but with the condensable volatile being methane instead of water. The climates of Titan, Triton, and Pluto, like those of Mars and Earth, are expected to vary in response to orbital forcings (e.g., obliquity and eccentricity), as volatiles are locked up in surface deposits and released in different epochs. Pluto’s orbit is so highly eccentric that its climate changes dramatically as it passes between perihelion and aphelion. In this section, the committee explores questions related to how planetary climates evolve.

Q6.2a What Processes Have Kept Earth’s Climate (Mostly) Clement over Geological Time, and What Has Been the Relative Role of Biological Versus Abiotic Feedbacks?

The Sun was ~30 percent less bright early in solar system history, yet the early Earth was not frozen. Evidence for surface liquid water dates back to almost 4.4 Ga (Valley et al. 2002). What greenhouse gases helped keep the surface warm? CO₂ and H₂O are thought to have been the two main candidates, but CH₄ could have played a role as well, especially prior to the rise of atmospheric O₂ (Haqq-Misra et al. 2008). CH₄ is largely biogenic, so this implies a degree of biological control of the evolution of Earth’s climate. A few percent of today’s CH₄ is produced by abiotic processes such as serpentinization. CO₂ is produced by volcanoes and is removed by weathering of silicate minerals, followed by deposition of carbonate sediments or veined carbonate in the seafloor. Its removal rate slows as the climate cools; thus, high CO₂ concentrations are an expected consequence of, and solution to, the faint young Sun problem (Walker et al. 1981). But what were early CO₂ levels and how rapidly did they decline with time? What geochemical proxies can we use to test this climate control hypothesis? Similar questions apply to atmospheric CO₂ and climate evolution on Venus and Mars, as discussed in Q6.2b,c.

Equally important to eukaryotic organisms, including higher plants and animals, was the evolution of atmospheric O₂. Prebiotic O₂ levels are believed to have been extremely low, of the order of 10⁻¹³ times the present atmospheric level, or PAL (Catling and Kasting 2017). Atmospheric O₂ remained low throughout the Archean, then rose abruptly during the Great Oxidation Event at ~2.4 Ga (Catling and Kasting 2017). Although the reasons why O₂ rose at this time continue to be debated, researchers agree that it was the evolution of oxygenic photosynthesis by cyanobacteria that ultimately led to this event. Cyanobacteria (formerly known as blue-green algae) are the only prokaryotic organisms that can perform oxygenic photosynthesis. The subsequent history of atmospheric O₂

during the ensuing Proterozoic and Phanerozoic eons is a research topic of great interest to astrobiologists because it is poorly understood and because it is relevant to the question of whether O₂ is a good biomarker on extrasolar planets (see Question 12, Chapter 15).

Q6.2b How Have Climate Conditions on Venus Evolved over Time, and Did the Planet Ever Have Surface Liquid Water in the Past?

Venus has a dense, 92-bar atmosphere consisting of 97 percent CO₂ but only 30 parts per million by volume of H₂O. The surface temperature on Venus is ~730 K, well above the critical temperature for water (647 K), so liquid water could not exist even if H₂O was abundant. Although Venus and Earth probably received an abundance of volatiles from planetesimals originating from farther out in the solar system (see Question 3, Chapter 6), Venus was closer to the Sun and experienced a higher solar flux, which could have triggered a runaway greenhouse (Ingersoll 1969). The steam atmosphere was photodissociated, hydrogen was lost to space, and the leftover oxygen either escaped along with the hydrogen or reacted with the crust. The buildup of CO₂ after that time was inevitable, as the weathering of silicate rocks to form carbonates was not possible once liquid water was no longer present (Urey 1952; Q6.1a). That said, the details of this process remain controversial. One hypothesis is that Venus developed a steam atmosphere during accretion and never had surface liquid water (Hamano et al. 2013). This model offers a straightforward explanation for how oxygen was lost: the greenhouse effect was so large that it kept Venus’s surface molten, allowing oxygen to react directly with a convecting magma ocean. This hypothesis is strengthened by recent 3D climate model calculations that support the idea that an initial steam atmosphere on Venus should never have condensed (Turbet et al. 2021). Both calculations conflict with a second hypothesis, also bolstered by 3D climate simulations, which suggest that if Venus started out as a slowly rotating planet with an ocean already present on its surface, dense cloud cover on the dayside could have kept surface temperatures well below the critical point of water (Way et al. 2016; Way and Del Genio 2020). In this latter scenario, liquid water could have remained on Venus’s surface for as long as 4 billion years until either increasing solar luminosity or some other planetary event such as enhanced volcanic outgassing destabilized it. Apparently, the assumed initial conditions can make a big difference in such calculations.

One way of distinguishing between these two competing scenarios is by determining the mineralogic composition of surface tesserae on Venus. Some investigators (e.g., Gilmore et al. 2015) have argued that the tesserae have a felsic composition that could only have been generated if liquid water was present. Evaluating this claim is a high priority for future Venus missions.

Q6.2c What Was the Nature of the Early Martian Climate, and How Were Conditions Allowing Rivers, Lakes, and Similar Features on the Surface Maintained?

While Mars today is cold and dry, overwhelming evidence indicates that large amounts of liquid water flowed over its surface for extended periods in the past (3–4 Ga). Mars’s more distant orbit makes this evidence even harder to explain than in the corresponding “faint young Sun problem” for Earth (Q6.2a). Various mechanisms have been proposed, including transient steam atmospheres from meteoroid impacts, sulfur-bearing gases from volcanism, or reducing gases such as hydrogen and methane from various sources, but all have challenges (Haberle et al. 2017). Understanding this period of martian history is essential for developing robust theories of exoplanet habitability (see Question 12, Chapter 15). The early evolution of Mars’s climate is also of key importance to planetary astrobiology because in situ analysis has shown that the planet was habitable to microbial life at this time (Grotzinger et al. 2014).

Current research on this problem is focused on understanding the timing and duration of these early warm episodes, and the nature of the water cycle at that time. Martian geochemistry provides a rich record to constrain early atmospheric composition (Ehlmann and Edwards 2014), but access to a range of samples from different environments is required. Isotopic analysis of rock and atmospheric samples returned to Earth, enabling far higher precision measurements than are possible in situ, will be particularly vital for constraining the chemical state of the early atmosphere. Some noble gas isotopes in the present-day atmosphere may be interpreted to infer the efficacy of early loss processes, while comparing elemental and isotopic composition of noble gases and light elements

between samples of ancient Mars (in returned rock samples and meteorites) and the present-day atmosphere provides information on atmospheric source and sink processes over time. Geochronological analysis of a range of surface samples is also needed to ascertain the timing of flowing water on Mars, and research into solar evolution is required to better constrain how the Sun’s bolometric and XUV luminosity has changed over geologic time. Last, further advances in atmospheric modeling are needed to understand the warming potential of various gases and aerosols and their impact on the observable rock record.

Q6.2d How Does Orbital Forcing, Including Obliquity and Eccentricity Changes, Govern Climate Change and Surface Volatile Redistribution on Extraterrestrial Bodies with Volatile Cycles Like Modern Mars, Triton, and Pluto?

Mars, Triton, and Pluto, and possibly large KBOs, have atmospheres controlled, or strongly influenced, by vapor pressure equilibrium—that is, the main atmospheric constituent (CO₂ on Mars, N₂ on Pluto and Triton) can partially or entirely condense on the surface; hence, the pressure of the atmosphere is regulated by the stability of surface ice temperature. In the case of Mars, CO₂ may also be exchanged with the regolith, greatly impacting pressure variations on orbital timescales (Buhler and Piqueux 2021). The temperature of surface ice varies primarily with insolation, which in turn varies seasonally but also with orbital parameters such as eccentricity and obliquity (Laskar et al. 2002). Changes in orbital parameters owing to Milankovitch cycles occur over tens of thousands of years and can impact the stability of ice formed from minor atmospheric constituents (e.g., H₂O on Mars, CH₄ on Triton and Pluto). Mars’s polar layered deposits are thought to record this orbital cycling of water ice and dust deposition, as dust abundance and transport also vary on orbital timescales. However, establishing a direct correlation is challenging because of uncertainties in how the water, CO₂, and dust cycles varied over time, the depositional/removal processes involved, and the detailed layer structure and composition (Byrne 2009). The presence of the vast N₂ ice deposit in Sputnik Planitia on Pluto stabilizes the pressure of Pluto’s atmosphere, although it is still subject to diurnal, seasonal and Milankovitch cycles (Bertrand et al. 2018). The distribution of ices on Triton remains largely unknown but could potentially provide information about volatile cycling in and out of the atmosphere. Large KBOs may have similar nitrogen atmospheres to Pluto and Triton and experience similar cycles over diurnal, seasonal and astronomical timescales (see also Q6.3b, Q6.4c,d).

Q6.2e What Is the Role of Transient Climate Forcing Owing to Meteoroid Impacts, Large Volcanic Eruptions, and Other Episodic Atmospheric Processes in Climate?

Transient climate forcing is important to many solar system objects. On Earth, large igneous province volcanic eruptions and meteoroid impacts are implicated in many past mass extinction events, while on Mars meteoroid impacts may have driven warming via direct thermal effects or alteration of the atmospheric composition (Q6.1). For Venus and Titan, it remains unclear if geologically recent transient outgassing events are responsible for their current climates, or if steady-state processes maintain the abundance of destructible species in their atmospheres. Changes in the rate of H₂O outgassing on Venus can have strong effects on albedo, and hence on climate, through coupling with the sulfur cycle (Bullock and Grinspoon 2001). The SO₂ atmosphere of Jupiter’s moon Io is supplied by volcanic eruptions discussed further in 6.4a. In all cases, it has become increasingly clear that transient processes play a major role in long-term climate evolution. Future progress in this area requires targeted geologic analyses to constrain the timing, duration, and volatile release of major volcanic and impact events, as well as laboratory experiments and numerical modeling to understand the physical and chemical behavior of atmospheres under the extreme conditions expected following a major volcanic eruption, meteoroid impact, or similar transient process.

Strategic Research for Q6.2

• Assess whether surface liquid water existed for an extended period on Venus by determining if the tesserae are felsic in composition.

• Constrain the timing of martian climate transitions by performing geochronological dating of samples from multiple locations on the planet’s surface.
• Constrain atmospheric evolution processes on Mars by returning samples of the atmosphere to Earth of sufficient concentration and fidelity to allow noble gas abundance and isotopic fractionation to be measured.
• Determine how and why Mars’s climate has changed over orbital timescales by performing radar and spectroscopic mapping of the polar layered terrain and by making in situ measurements of their structure and composition (thickness of layers, dust content, and isotope ratios) and their local meteorology (including volatile and dust fluxes).
• Study the annual climate cycles and long-term evolution of atmospheres controlled by vapor pressure equilibrium by measuring the isotopic composition and fractionation of atmospheric gases and, if possible, of surface ices and atmospheric haze, on Triton and Pluto.
• Study surface-exchange processes across a diverse range of atmospheric compositions by developing one-, two-, and three-dimensional models of past and present planetary climates.
• Determine accurate models of the extreme ultraviolet solar luminosity evolution over geologic time by observing a variety of solar-type stars using space-based observatories.
• Investigate the radiative forcing potential, chemistry and microphysics of greenhouse gas, haze, and cloud combinations relevant to climate evolution processes by performing laboratory studies.

Q6.3 WHAT PROCESSES DRIVE THE DYNAMICS AND ENERGETICS OF ATMOSPHERES ON SOLID BODIES?

The dynamics and energetics of solid-surface atmospheres are largely driven by the distribution of solar radiation, which is governed by the solid body’s orbital parameters. Atmospheric composition, surface properties, and magnetic environment also impact how solar radiation forces atmospheric circulations. Many outstanding questions relating to atmospheric dynamics and energetics remain unanswered (see Figure 9-3).

Large-scale, meridional (latitude-height) overturning circulations, such as Hadley cells, transport heat, momentum, and tracers between different atmospheric regions. Surface topography, boundary-layer convection, and daily patterns of surface heating also produce various atmospheric waves, which can travel large distances before breaking and depositing their energy and momentum, modifying the angular momentum structure of an atmosphere. Waves are thought to be responsible for strong atmospheric superrotation and for large perturbations to the dynamics of upper atmospheres, but the proposed mechanisms remain mostly unexplored.

The surface-atmosphere transfer of heat, momentum, gases, volatiles, and dust, including their mixing within the lowest portion of the atmosphere (the planetary boundary layer), has a major impact on much of the observed weather and climate, yet has rarely been measured other than on Earth. This is especially problematic when such processes dominate the atmospheric circulation and its variability, as is the case for lofted dust on Mars or condensation-sublimation flows on Pluto, Triton, and Io. Aerosols (dust, hazes, and clouds) are key features of many planetary atmospheres and affect the atmospheric absorption and scattering of solar radiation. Aerosol transport, microphysics, and radiative processes are often coupled via complex feedbacks, yet few measurements related to such processes exist for atmospheres other than Earth’s.

At higher altitudes, an ionosphere is generated by the interaction between high-energy particles from the Sun and the thermosphere. The ionosphere and thermosphere are thus coupled chemically and dynamically via energy and momentum transfer (ion-neutral drag). At still-higher altitudes, the ionosphere is coupled to the magnetosphere and through it to the surrounding space plasma. Magnetospheres are driven from above by the highly variable solar wind and from below by particles escaping from the planetary atmosphere and ionosphere, with these interactions varying if a planetary magnetic field exists. These upper-atmosphere couplings and interactions have never been fully characterized for any planetary body, thus their full impact remains poorly known.

Surface-bounded exospheres have their own unique dynamics based on the specific atomic and molecular species present, surface temperatures, surface solar-wind and meteoroid exposure, and the loss processes for the species. The dynamic formation and migration of exospheric volatiles may have been critical in forming the polar volatile deposits at the Moon, Mercury, and Ceres (see also Q5.5c). However, it remains unclear whether some

volatiles, such as water, can efficiently migrate to polar cold traps and/or have a surface-exosphere cycle that operates on short timescales.

Studying present-day atmospheric circulations, surface fluxes, and solar inputs on bodies with differing orbital settings, atmospheric compositions, and surface properties thus provides valuable insight into a huge range of fundamental dynamical processes. This helps us to not only understand the atmospheres of present-day Earth and other solid bodies in the solar system, but also to extrapolate this knowledge to the past climate states of those bodies (see Q6.1, 6.2) and to the atmospheres of rocky exoplanets (see Q12.6).

Q6.3a How Do Horizontally and Vertically Propagating Waves Drive Planetary Atmosphere Dynamics?

The atmospheres of the solar system’s slowest rotators, Venus and Titan, exhibit strong equatorial superrotation, with weaker equatorial jets also found on the more rapidly rotating Earth and Mars (Read and Lebonnois 2018). Based on dynamical arguments, equatorial superrotation requires the wave-driven transport of energy up the

angular-momentum gradient. Yet while mechanisms for driving this have been proposed, most of the atmospheric waves involved have not been observed. Furthermore, it is unclear whether Venus’s upper atmosphere superrotation is driven by the same mechanism as the observed superrotation at the cloud tops and below.

Recurring orographic gravity waves above venusian mountain ranges substantially torque the solid body (Navarro et al. 2018), while possibly reaching Venus’s upper atmosphere and may in turn impact atmospheric dynamics at high altitudes. Even less information is available for vertical wave propagation on Titan. On Mars, increased thermal tides and convectively driven gravity waves during dust storms likely help to drive water to higher altitudes and result in greater atmospheric loss (Yiğit et al. 2021). The influence of lower atmosphere waves on thermospheric and ionospheric dynamics is beginning to be examined in models, but more data are needed to test these predictions.

Crucially, the episodic and often unpredictable nature of wave-driven dynamical changes can only be understood via long-term atmospheric measurements of planetary atmospheres, combined with theoretical/modeling studies.

Q6.3b What Controls the Onset, Evolution, and Year-to-Year Variability of Dust Storms on Mars and Titan?

Dust clouds have long been observed in the martian atmosphere, and more recently in Titan’s. While the impact on Titan’s thick atmosphere is unknown, regional and global dust storms have a huge impact on radiative heating in the thin martian atmosphere, and hence on temperature and winds (e.g., Kahre et al. 2017), as well as on atmospheric loss rates (see Q6.5b). The trigger for global storms, which shroud Mars in dust and completely change the atmospheric circulation for several months, is not currently understood. Modeling suggests that surface-dust availability, water-ice nucleation on dust particles, and complex feedbacks between dust lifting, transport, and circulation patterns (including waves) may all be important (e.g., Newman et al. 2016). Fundamental unknowns include how and where dust is lifted from the surface (e.g., whether direct lifting or saltation of sand particles is involved), how that dust is then raised through the boundary layer and free atmosphere, and how the dust and water cycles are coupled. The infrequency (typically three per 20 Earth years) and significant variability of global storms further complicate understanding their origins and motivates long-term measurements of Mars. Additionally, the presence of high-altitude dust layers outside of dust storms remains largely a mystery (e.g., Heavens et al. 2014).

Q6.3c How Do Sublimation–Condensation Flows Drive Circulation in Thin, Transient Atmospheres? What Is the Nature of This Circulation and How Does It Relate to Surface-Ice Distributions?

Far from the Sun, the surface pressures of Pluto and Triton vary by orders of magnitude over their long orbits as their N₂ atmospheres largely condense out onto the surface around aphelion. Modeling suggests their atmospheric circulations are dominated by surface topography and the distribution of surface ice, which control both thermally driven and sublimation-condensation driven flows (e.g., Bertrand et al. 2018).

On Io, the contribution of volcanic emission versus sublimation of SO₂ surface ices to the varying surface pressure, which varies spatiotemporally over more than five orders of magnitude, is vigorously debated (e.g., Tsang et al. 2016). Day-night, volcanic, and eclipse-driven temperature differences cause huge pressure differentials and fast, complex circulations, complicated further by plumes (McDoniel et al. 2017). On Triton, such phenomena may also be important. Constraints on these models, however, are very limited as the distribution of surface ices is basically unknown for Triton and Io, and only partially known for Pluto, and the predicted volatile fluxes and complex circulations are unmeasured (see also Q6.4d).

Q6.3d How Do Haze and Cloud Processes Impact Atmospheric Dynamics, and What Are the Key Couplings and Feedbacks?

Haze layers impact radiative heating via absorption and scattering, and are present in at least Venus, Titan, Pluto, and Triton’s atmospheres. Models suggest that haze-dynamical feedback may be an important driver of Titan’s stratospheric circulation, but this is unclear from observations. On Venus, unknown ultraviolet absorbers

in the thick clouds produce significant albedo variations (hence amount of solar energy absorbed) on decennial timescales, potentially modifying superrotation (Pérez-Hoyos et al. 2018). However, despite decades of study, the nature of these absorbers remains unknown, as does their impact on dynamics. On Titan, latent heat release and strong convection are associated with methane clouds. The processes by which hydrocarbons form clouds (Q6.6) and evaporate from Titan’s surface seas and lakes (Q6.4)—analogous to those involved in Earth’s water cycle—have never been observed in situ and are unlikely to be encountered by the Dragonfly mission, which will target the dry low latitudes during a period expected to have minimal precipitation.

Q6.3e What Determines the Effectiveness of Ion-Neutral Drag on Augmenting Upper Atmospheric Circulation?

For Venus, Mars, and Titan, the transfer of energy and momentum between the neutrals and the ionospheric plasma is not well quantified. The measurement of upper-atmosphere neutral and ionospheric winds by current and past missions has been incomplete and inconsistent. At Venus, polar asymmetric ionospheric ion flow has been observed, but it is unknown if this asymmetry drives upper-atmosphere superrotation through ion-neutral drag (Lundin et al. 2011). Moreover, the interaction of the neutral gas and ionosphere in the southern hemisphere of Mars is mediated by strong remanent crustal magnetic fields, but the dynamical effects of these fields for both the neutrals and ions is poorly understood. Understanding the ionosphere structure, variation, and drivers feeds back into knowing how it connects to the neutral atmosphere.

Q6.3f How Do the Structure and Dynamics of Planetary Magnetospheres Vary with Season and Solar Inputs?

Planetary magnetospheres represent a transitional region between the upper atmosphere and the interplanetary-space environment (Russell 2001; Kivelson and Bagenal 2014). The nature of the interface depends critically on the properties of the solar wind and the interplanetary magnetic field (IMF) that it carries. For intrinsic magnetospheres, such as those of Earth and Mercury, the IMF orientation controls the flow of matter and energy from the solar wind to the magnetosphere by magnetic reconnection and boundary layer processes such as the Kelvin-Helmholtz instability. For induced magnetospheres, such as those of Mars and Venus, the IMF determines the orientation of the entire magnetosphere. In either case, the magnetosphere is driven both from above by the highly variable solar wind energy, momentum, and electromagnetic fields, and from below by escaping particles from the planetary atmosphere (see Q6.5c). Magnetospheric structure and dynamics thus vary both with season and solar cycle. The relative importance of these competing influences depends on a wide range of parameters and leads to a complex and variable interaction that has yet to be well characterized for most solar system objects.

Q6.3g What Controls the Transport and Sequestration of Volatiles in Solid-Surface Exospheres?

Molecular hydrogen and ⁴⁰Ar in the Moon’s atmosphere may undergo numerous surface adsorption-desorption sequences to effectively migrate across the body and sequester in cold regions, such as the lunar nightside and polar cold traps. Sequestered nightside volatiles can be rereleased via thermal desorption when the surface rotates into daylight, forming a volatile cycle. However, it is unclear if released water molecules, an important lunar trace species, undergo this same migration and cycling. Early modeling suggested that solar-wind protons convert to regolith water, which is eventually released and migrates to polar cold traps, implying that the cold traps are currently active and dynamic. However, the lack of a detectable stable water exosphere on the Moon suggests that released water may be dissociated/destroyed at the surface after a single adsorption-desorption hop, preventing water from migrating to the poles (Benna et al. 2019). This raises the question of whether water and other volatiles (such as CO, CO₂, and methane) can migrate on exposed rocky bodies such as the Moon, Mercury, Ceres, and Phobos, and whether they form dynamic cycles or are lost to the surface immediately upon release. Recent research suggests water may be released from Mercury’s regolith and may migrate to the poles (Jones et al. 2020),

but it remains unclear if surface formation/release, transport, and sequestration/trapping of volatiles are universal processes across all exosphere-only rocky bodies.

Strategic Research for Q6.3

• Determine how atmospheric waves drive atmospheric dynamics and energetics, especially phenomena such as supperrotation and lower-upper atmosphere coupling by observing wave amplitudes, periods, phases, and spatiotemporal distributions in thermal and direct wind measurements over multiple annual cycles on Venus, Mars, and Titan.
• Determine how the surface is coupled to the main atmosphere by measuring the transport of heat, momentum, volatiles, and dust through the planetary boundary layer, via in situ and remote sensing observations of fluxes covering key time periods or environmental conditions, on bodies with collisional atmospheres such as Venus, Mars, and Titan.
• Investigate the cause of variability in martian dust storms and hence climate by making in situ measurements of surface dust and sand fluxes simultaneous with environmental conditions, and in situ and orbital measurements of surface dust and sand availability.
• Determine how the atmospheric circulation is driven on bodies with thin, transient atmospheres by measuring the thermal state and winds, and distribution of surface topography, ices, and (where relevant) plumes, via remote sensing of Pluto, Triton, and Io.
• Determine how aerosols influence atmospheric dynamics and energetics by measuring their properties and spatiotemporal distributions, simultaneous with the thermal and circulation response of the atmosphere, on diurnal, seasonal, and multi-annual timescales on Venus (clouds and hazes), Mars (dust and clouds), and/or Titan (hazes, clouds, and dust).
• Determine the effectiveness of ion-neutral drag on augmenting upper atmospheric circulation by performing in situ measurements of ion and neutral winds, as well as ion electron densities, plasma distribution functions, and magnetic fields, on Venus, Mars, and Titan.
• Determine how the magnetospheres of solid bodies are driven by solar inputs from plasma timescales to solar cycle timescales by simultaneously measuring the upstream solar wind input and the magnetospheric response at Mercury, Venus, and Ganymede.
• Investigate the mechanisms by which waves are generated and interact with the mean flow in planetary atmospheres to better understand phenomena including equatorial superrotation (e.g., Venus and Triton), lower-upper atmosphere coupling, atmospheric loss, and martian dust storms, using numerical models.
• Explore the connections between the solar wind, magnetic fields, and the neutral atmosphere/exosphere through numerical modeling.

Q6.4 HOW DO PLANETARY SURFACES AND INTERIORS INFLUENCE AND INTERACT WITH THEIR HOST ATMOSPHERES?

Volcanism and other forms of interior outgassing have played a pivotal role in the formation of atmospheres on various bodies, including Earth, Mars, Venus, Io, and Titan (see Q6.1), and have continued to influence those atmospheres throughout their history. Io, for example, is the most volcanically active world in the solar system, and near-continuous volcanic outgassing is the source of both its SO₂ atmosphere and a global sulfuric acid cloud layer. On other bodies, including Mars, Venus, and Titan, the importance of modern outgassing remains a mystery. While some observations point to modern and active volcanism on these worlds, definitive evidence for these processes is lacking (Q6.4a). In addition to volcanic processes, plumes can deliver particulates that affect the atmospheric chemistry, composition, and dynamics of bodies such as Europa, Enceladus, and Triton (Q6.4b). On Mars, CO₂ jets loft dust into the atmosphere each spring, but the importance and detailed mechanics of this source remain unknown (Q6.3b). The extent to which Triton’s plumes influence atmospheric composition and the moon’s haze layers similarly remains a mystery. Images of several of these bodies are shown in Figure 9-4.

Surface-atmosphere interactions encompass a multitude of processes that link the surfaces of planetary bodies to their atmospheric boundary layers or surface-bounded exospheres. Interactions are bi-directional: they include energy and material exchanges (e.g., condensation and sublimation) and can also include feedback mechanisms that amplify or attenuate coupled processes (e.g., positive feedbacks between dust lifting and circulation strength).

Observations of surface changes caused by aeolian processes often constitute our only information on near-surface wind patterns and environmental conditions. The composition of volatile species in planetary atmospheres drives the formation, evolution, and stability of ices across planetary surfaces. The mechanisms by which condensable volatile reservoirs drive material transport, govern the evolution of large deposits including polar caps, and generate clouds and weather within their host atmospheres remain poorly understood.

In addition to Earth, Titan and early Mars are known to support, or have supported, active hydrologic systems that incorporate stable bodies of surface liquid. Together, they provide natural laboratories for investigating how hydrologic cycles operate and exchange material between atmospheric, surface, and subsurface reservoirs over diurnal, seasonal, and orbital timescales. On airless bodies such as the Moon, Mercury, and Ceres, polar cold traps are of compelling interest to investigations of how volatiles inform the processes of prebiotic chemistry, as well as the history of volatiles throughout the age of the solar system. Our understanding of the volatile/organic composition, evolution, and astrobiological potential of the deposits in these regions, however, remains incomplete. For example, despite observations that surface ice is sparse in lunar polar cold traps (Hayne et al. 2015), it is believed that the ice in these cold traps should be in sublimation–condensation vapor-pressure equilibrium (Zhang and Paige 2009). Understanding this will require comparing the structure, content, and concentration of the trapped volatiles on the Moon to those at Mercury and Ceres (see also Q5.5c).

Q6.4a How Does Modern Volcanic Outgassing from a Planet’s Interior Influence the Composition of Its Current Atmosphere?

Following initial volcanic outgassing, SO₂ on Io subsequently experiences diurnal condensation and sublimation. The dynamics of SO₂ exchange processes, including supersonic flow from the dayside to the nightside, is not well understood. On Mars, Venus, and Titan, it is unknown how important modern outgassing might be. The youngest lava flows and meteoritic evidence on Mars suggest it was active as recently as ~50,000 years ago (Horvath et al. 2021). The InSight lander has measured hundreds of “marsquakes” and continues to monitor seismic activity that may be related to active volcanism. There is circumstantial evidence for ongoing volcanic activity on Venus, including transient thermal emissions from areas of stratigraphically young volcanic deposits and residual heat emitting from young lava flows that lack evidence of weathering. Decadal changes in atmospheric composition suggest that SO₂ may be reinjected into Venus’s upper atmosphere by major volcanic eruptions, and the presence of a global sulfuric acid cloud is, as on Io, thought to be maintained by continuous volcanic outgassing. While Doom Mons and Sotra Patera on Titan have been interpreted as a cryovolcano and cryovolcanic caldera, respectively (Lopes et al. 2013), individual flows have not been identified. Regardless, modern outgassing is one hypothesis for the replenishment of methane in Titan’s atmosphere.

Q6.4b How Do the Particulates Delivered by Plumes Affect the Atmospheric Chemistry, Composition, and Dynamics of Their Host Bodies?

Every spring, cold jets of CO₂ sublimating from seasonal polar caps loft dust into the martian atmosphere (Kieffer et al. 2006). Understanding how these plumes affect the martian climate will require long-duration monitoring observations that encompass multiple dust storm events. At least 2 and possibly as many as 14 plumes erupting on Triton were observed by Voyager 2, along with more than 120 dark streaks and fan features thought to be the remains of plumes (Smith et al. 1989). While the mass flux of these plumes was calculated, other basic properties such as composition remain unknown. The origin of the plumes is similarly unknown, with hypotheses ranging from surficial N₂ sublimation to eruption of subsurface water. Their influence on Titan’s atmosphere is also unknown. Understanding plume composition can reveal fundamental properties that will reveal their source and constrain their influence on Triton’s atmosphere.

While evidence supports the possible eruption of water vapor plumes from Europa, we know very little about them if they do exist. Are particulates present? If so, what is their composition? Once deposited, do they leave an identifiable mark on the surface, such as color or photometric properties?

Q6.4c What Can Aeolian Features and Their Temporal Variation Tell Us About Near-Surface Winds and Other Atmospheric Variables? What Are the Mechanisms That Connect Wind-Driven Surface Modifications to Environmental Conditions?

Aeolian processes are important for lofting dust into the atmosphere of Mars and for modifying its circulation (see Q6.3b). Dust clouds have also been observed on Titan and possibly Venus. Aeolian processes are also vital for shaping and reshaping planetary surfaces by modifying albedo, forming, and migrating dunes, sculpting rocks, and more (see Question 5, Chapter 8), with such features observed on bodies as diverse as Venus, Pluto, and even comet 67P Churyumov-Gerasimenko. While aeolian-driven surface changes typically have only minimal impact on the state of the atmosphere, observing such features is often the only way by which the near-surface circulation of a planetary atmosphere can be inferred. Therefore, understanding the mechanisms that connect atmospheric variables (such as wind stress and direction) to aeolian features on bodies other than Earth is vital for extracting information on the modern near-surface atmosphere, while ancient aeolian features provide insight into the past. Major unknowns remain, such as how particulate lofting and transport varies with material composition and environmental factors (e.g., gravity and atmospheric density), along with the processes that lead to dune formation in varying environmental conditions.

Q6.4d How Do the Major Constituents in Planetary Atmospheres Drive the Formation, Evolution, and Stability of Polar Caps and Other Reservoirs over Seasonal Timescales? How Do These Reservoirs Drive Volatile Transport and Generate Winds That Can Result in Aeolian Activity?

The sublimation and condensation of volatiles in atmospheres controlled, or strongly influenced, by vapor pressure equilibrium (e.g., Mars, Triton, and Pluto) (e.g., Leighton and Murray 1966) are driven predominantly by energy balance. There are nuances (such as albedo of frost versus ice), however, that are not well understood. These factors can affect the rates of sublimation and condensation that, in turn, drive winds as the volatiles move between cold traps (e.g., polar caps, Sputnik Planitia; see also Q6.3b). The nature of the surface topography also affects local energy balance. These subtleties can have profound effects on polar cap boundaries, weather, and climate. For example, regional dust storms on Mars affect the amount of spring sublimation activity, while spring sublimation drives the initiation of local dust storms along the edge of the retreating seasonal cap and winds associated with cap thermal contrasts may be key to generating regional and even global dust storms (see also Q6.3b). Does the level of dust in the atmosphere affect the amount of snowfall, which has been observed to vary from year to year? Over orbital timescales encompassing obliquity and eccentricity cycles, how does the energy balance change, and are polar cap boundaries affected?

Q6.4e How Do the Minor Constituents in Planetary Atmospheres Drive the Distribution of Surface and Near-Surface Volatiles, Such as the Distribution of H₂O Ice on Mars or CH₄ and CO Ices on Triton/Pluto?

For atmospheres dominantly controlled by sublimation/condensation equilibrium, volatile ices on the surface are also present as gases in the atmosphere. Depending on the nature of surface-atmosphere interactions, however, the abundances of these species can vary by orders of magnitude. The abundance of minor constituents influences atmospheric processes, including cloud generation and weather, and determines the redistribution of surface volatiles over seasonal and orbital timescales. Unknowns include how stable water ice deposits form and survive on Mars, as well as the distribution and migration of volatile methane and carbon monoxide ices on bodies with tenuous atmospheres like Triton and Pluto. Over orbital timescales, the role of minor constituents also changes. Understanding the influence of water ice and dust cycles in the polar layered deposits in response to variations in, for example, obliquity and eccentricity, analogous to Milankovitch cycles on Earth, is important for interpreting how climate differed over orbital timescales (Q6.2d). The layering seen in the martian polar caps is assumed to reflect such changes; however, it is challenging to identify a direct correspondence (Becerra et al. 2019). To extrapolate back in time, the microphysics of the current climate needs to be better understood. For example, how much ice is condensed every year, how much is sublimed, and what sort of “dust lag” (if any) is left behind?

Q6.4f What Do the Present-Day Methane Cycle on Titan and Past Water Cycle(s) on Mars Tell Us About How Hydrologic Cycles Operate and Exchange Material Between Surface and Atmospheric Reservoirs?

Titan is the only solar system body other than Earth to have stable surface liquid and major latent heat exchange within its hydrologic cycle, with methane and ethane taking the place of water. As a result, Titan’s methane-based hydrologic cycle is an extreme analogue to Earth’s water cycle and represents a natural laboratory to study how planetary climate and hydrologic cycles operate and are maintained (Hayes et al. 2018). Many aspects of Titan’s hydrologic system remain unknown, including the composition and full seasonal variability of clouds, the importance of surface-atmosphere exchange and surface/subsurface transport, and the importance of ethane. For example, clouds were expected to form on Titan at northern mid-latitudes in the period following spring equinox, but their appearance was delayed for unknown reasons (Turtle et al. 2018). Better identification of cloud formation mechanisms, their links to surface sources, and whether clouds are composed of ice, liquid, or a binary methane-nitrogen solution, are required to understand the present climate and interpret surface evidence for paleo-seas and fluvial activity. Titan’s present hydrologic cycle is also a natural laboratory for investigating exchange processes between atmospheric, surface, and subsurface reservoirs over multiple timescales. Evidence suggests that Mars also had a water-based hydrological cycle in the past (Q6.2c), but it is unclear how the location and size of a possible ocean would have affected the nature of this hydrological cycle, especially during periods when other surface liquid was marginal or short-lived, or how water ice clouds would have affected the radiation balance in a warmer and wetter climate.

Strategic Research for Q6.4

• Determine if modern volcanic or tectonic activity is influencing Venus’s atmosphere using thermal emission, infrasound, and/or ultraviolet emission from breaking atmospheric waves.
• Investigate the impact of volcanic outgassing on Io’s atmospheric composition by documenting volcanic activity on Io over multiple seasonal cycles using ground- or space-based observatories.
• Investigate the redistribution and transport of liquid and solid hydrocarbons on Titan over time by mapping the distribution of geologic features, volatile ices, and tropospheric clouds over multiple Titan seasons, utilizing ground- and space-based observatories to provide a more extended temporal dataset.
• Investigate the source of Triton’s plumes and their contribution to its atmosphere by measuring the composition of plume material.
• Study the properties of Triton’s vapor-pressure atmosphere, including why/how it is different from Pluto, by measuring the distribution of surface ices, as well as atmospheric pressure and temperature as a function of altitude.
• Investigate how and where stable water ice deposits form on Mars by measuring their distribution through radar and spectroscopic mapping from orbit, and by measuring the ice vertical distribution, volatile fluxes, and environmental drivers at the surface.
• Determine how dust lifting and sand motion are linked to the state of the near-surface atmosphere by making simultaneous, in situ measurements of dust and sand fluxes, surface properties, and environmental conditions (e.g., winds and electric fields) at the surface of Mars, Venus, and Titan.
• Test and improve theories relating dust lifting, sand motion, and aeolian feature characteristics to environmental conditions (e.g., winds, electric fields, grain sizes) on other planetary surfaces by performing laboratory, numerical, and terrestrial analog field studies.
• Infer near-surface wind patterns and environmental conditions, and how they change over time, by documenting aeolian features and how they vary over seasonal and annual cycles via high-resolution imaging on Mars and Titan.

Q6.5 WHAT PROCESSES GOVERN ATMOSPHERIC LOSS TO SPACE?

Atmospheres and exospheres of every variety escape to space through a richly varied set of processes (Figure 9-5). Atmospheric escape is inextricably linked to evolution (Q6.2), with loss processes coupled to interiors and surfaces (Q6.4), as well as to the Sun and the space environment and to impacts in the early solar system (Q6.1).

Escape can alter the chemistry of an atmosphere by preferentially removing some species or isotopes (typically the less massive constituents), or it can transform an early thick atmosphere to a thin atmosphere, as is thought to have occurred on Mars (Jakosky et al. 2018). The net escape rate depends on both the reservoir of particles and the external energy inputs, and thus may be either supply-limited or energy-limited. Atmospheric escape pathways include both thermal and nonthermal processes (Lammer et al. 2008; Gronoff et al. 2020). The relative importance of thermal mechanisms such as hydrodynamic and Jeans escape depends on object size (through gravity) and upper atmospheric temperature, and thus on the evolutionary stage of the body and the Sun (see Q6.1b). Many nonthermal processes involve both neutral and charged particles, thereby coupling atmospheric escape to the ambient plasma environment.

The presence of a magnetic field alters the interaction with the solar wind (see Figure 9-5), reducing the efficiency of some escape processes, but increasing the efficiency of others; however, the net effect of a magnetic field on atmospheric escape remains unquantified. Although all escaping particles ultimately pass through the collisionless exosphere, the structure, composition, chemistry, and dynamics of the collisional atmosphere constrain what species can travel from the surface to the exosphere and escape. Escaping particles from planets and their satellites interact with the solar wind and other plasmas to drive the structure of planetary magnetospheres, in turn affecting the escape pathways.

Surface boundary exospheres involve their own unique set of processes (Dukes and Hurley 2016). In contrast to the collisional case, the space environment directly interfaces with the surface of these bodies, influencing both the delivery and loss of volatiles (Figure 9-6). Magnetic fields, both global and local, impact plasma-driven delivery to and loss of volatiles from the surface. To fully understand surface boundary exospheres, their equilibrium states, and their long-term evolution, we need a better understanding of the surface physics and chemistry and quantify the relative contributions from all relevant energy sources, including their seasonal and solar cycle variations. Atmospheric loss discussed herein also applies to those bodies having quasi-collisional or transient atmospheres, temporarily evolving from a collisionless exosphere to a collisional atmosphere and back. The early Moon and large KBOs are possible examples of such systems. Additional remarks on escape from quasi-collisional atmospheres can be found in Q6.1b.

Q6.5a How Does the Presence or Absence of Intrinsic or Parent Body Magnetic Fields Influence the Escape of Gases from Solid Planets and Satellites?

The hypothesis that a global magnetic field shields the atmosphere from solar wind-driven loss processes seems intuitively obvious, as the solar wind plasma and its electromagnetic fields cannot then interact directly with the upper atmosphere. However, a global field also increases the interaction cross section of a planet, causing it to intercept more solar wind energy that can drive outflow through the polar cap and cusps, competing with shielding effects and in some scenarios even enhancing loss (Gunell et al. 2018; Gronoff et al. 2020). Similarly, parent body magnetic fields can channel magnetospheric plasma to satellites (see Q8.4), potentially enhancing atmospheric loss. Among our limited sample of terrestrial planets, Venus has a thick atmosphere but no global magnetic field, Earth has a global field and a substantial atmosphere, and Mars has no present-day magnetic field and a much thinner atmosphere. Yet all three have comparable charged particle escape rates. In the outer solar system, Io, Titan, and Triton are located within giant planet magnetospheres and thus are subject to very different incident plasma than Pluto in the solar wind. Although some constraints on thermal escape exist for Pluto and Titan, the nonthermal escape rates and drivers of escape remain poorly constrained for all four bodies. Ganymede presents a unique scenario, with a significant intrinsic magnetic field embedded in the intense jovian field. While the moon’s intrinsic field could potentially shield its atmosphere, reconnection with the parent body field may enable escape. Ultimately, while the planets and satellites in the solar system provide a variety of endmembers, they have so many confounding differences (notably size and atmospheric composition) that the importance of magnetic fields remains uncertain. Thus, we do not yet know whether an Earth-like atmosphere can accumulate and remain stable without a global magnetic field, or whether the cessation of the martian dynamo early in its history played a significant role in the loss of its early atmosphere (Lillis et al. 2008).

Q6.5b How Do Atmospheric Dynamics, Such as Martian Dust Storms, Affect the Escape of Gases from Solid Planets and Satellites?

While atmospheric gases, by definition, escape through the exosphere, processes that occur below the exobase can still have significant effects on loss rates. At Mars, although diffusion-limited transport of H₂ would lead to only slow variations in hydrogen loss, we now know that hydrogen escape rates can vary by up to an order of magnitude during dust storms (see also Q6.3b). The observed correlation between dust storms, increased high-altitude water content in the upper atmosphere, and enhanced hydrogen escape suggests that coupling between the lower and upper atmosphere may enable the transport of water to high altitudes, thereby driving more rapid variations in hydrogen escape (Heavens et al. 2018). Atmospheric dynamics may facilitate escape directly and/or by coupling to ionospheric and magnetospheric processes (see Q6.3e,f), with the latter particularly relevant for unmagnetized bodies, such as Venus, Mars, and Pluto (see Q6.5c). Upper atmospheric processes also influence escape through photochemistry (see Q6.6d), with dissociative recombination capable of providing particles with escape energy at less massive bodies, such as Mars, Titan, Triton, and Pluto. The complex coupling from the lower atmosphere to the upper atmosphere to the exosphere implies that physics in multiple regions may affect the total escape, but in many cases the limiting processes remain unknown. To advance our understanding of atmospheric escape, we therefore need to make more complete measurements of the energetic inputs and develop coupled models capable of capturing the physics in the different regions of the atmosphere.

Q6.5c How Do Escaping Gases Influence the Structure of Planetary Magnetospheres?

Escaping neutral and charged particles influence the structure of planetary magnetospheres. Escape from planetary satellites provides an important (even primary, for some giant planets) source of mass and momentum in their parent body magnetospheres (Q8.4). At bodies with induced magnetospheres, such as Venus, Mars, Titan (when in the solar wind), Pluto, and even comets, the upper atmosphere interfaces directly with the space environment, providing the primary obstacle to the solar wind (Luhmann et al. 2004). However, the relative importance of induced currents that produce a magnetic barrier and thus an obstacle to the solar wind flow, as opposed to

mass loading in the extended exosphere that transfers momentum from the upstream solar wind, remains poorly quantified and may vary significantly with, for example, object size and solar inputs. For bodies with intrinsic magnetic fields, such as Earth (Toledo-Redondo et al. 2021) and Mercury, plasma derived from the atmosphere and/or exosphere can also affect the magnetosphere by influencing large-scale current systems and altering local dynamics (e.g., magnetic reconnection at the magnetopause). However, the processes by which escaping gases influence magnetospheric structure and dynamics, and the details of the feedback between solar wind energy input, atmospheric escape, and magnetospheric structure remain incompletely understood for both intrinsic and induced magnetospheres.

Q6.5d How Do Magnetic Fields Influence the Loss of Volatiles from Objects with Surface Boundary Exospheres?

The presence of local and/or global magnetic fields can also affect the escape of volatiles from objects with surface boundary exospheres. By shielding portions of the surface from charged particle bombardment, magnetic fields can reduce the release of volatiles, thereby decreasing the net loss rate by direct escape to space and/or subsequent loss processes (Poppe et al. 2014). Magnetic fields can also regulate the loss of charged particles by shielding them from the ambient plasma and preventing them from being picked up. At Mercury, some portions of the surface routinely experience plasma bombardment, with others only exposed during more extreme solar wind conditions. At the Moon, localized portions of the surface have significant magnetic fields capable of deflecting or reflecting the solar wind. At Ganymede, the polar and equatorial portions of the surface interact with very different incident charged particle populations (as demonstrated by surface color variations). What are the local and global effects on escape from these bodies?

Q6.5e How Is the Escape of Volatiles from the Moon, Mercury, and Other Bodies with Surface Boundary Exospheres Driven by Photon, Charged Particle, and Micrometeorite Influx?

The escape of volatiles from bodies with surface boundary exospheres is driven by a variety of energetic inputs (Dukes and Hurley 2016). Photons stimulated desorption and charged particle sputtering remove particles from the surface in both neutral and charged form, and micrometeoroid impacts also liberate material. Solar photon irradiation stimulates thermal desorption of atoms and molecules from the surface. Photons can also facilitate the escape of some exospheric species by exerting radiation pressure, and by ionizing and/or dissociating neutral particles. Once ionized, exospheric constituents can be picked up and accelerated to escape energy and beyond by the ambient plasma, via electromagnetic fields. These processes occur to some degree at all surface boundary exospheres, but our understanding of their relative importance, their similarities and differences, and their dependence on parameters such as surface binding energy, desorption activation energy, solar activity, body mass (gravity), and distance from the Sun, is just beginning to take shape (Killen et al. 2018). We have not yet quantified the relative importance of these escape processes in comparison to sequestration in the regolith or in cold traps, nor understood the variations in the efficiency of these processes with season and solar cycle.

Strategic Research for Q6.5

• Trace the flow of energy and escaping gases through collisional atmospheres (e.g., Venus, Mars, Triton, Titan, Pluto) to diagnose lower–upper atmosphere coupling, utilizing simultaneous measurements of the lower and upper atmosphere and exosphere.
• Discover how escaping ions influence the magnetospheric current systems by performing multi-point measurements in induced (e.g., Venus and Mars) and intrinsic (e.g., Mercury) magnetospheres.
• Relate the loss of volatiles from surface boundary exospheres to solar wind dynamics and quantify the effects of magnetic fields by measuring escaping, migrating, and bound species in regions of different magnetic topology (e.g., Mercury and Ganymede polar and equatorial regions and lunar magnetic anomalies).

• Reveal the factors that control the structure, composition, and dynamics of surface boundary exospheres (e.g., Mercury, Moon, Ceres, and Europa) by simultaneously measuring energetic inputs and escaping species for at least one orbit and preferably for a substantial portion of the solar cycle.
• Establish the influence of magnetic fields in constraining atmospheric loss by simulating escape from bodies with and without intrinsic and/or parent body magnetic fields with sufficient fidelity to resolve pertinent loss processes.
• Investigate how photons, charged particles, and micrometeoroids drive escape from different types of surfaces by performing desorption, sputtering, and impact vaporization (i.e., space weathering) laboratory experiments.

Q6.6 WHAT CHEMICAL AND MICROPHYSICAL PROCESSES GOVERN THE CLOUDS, HAZES, CHEMISTRY, AND TRACE GAS COMPOSITION OF SOLID BODY ATMOSPHERES?

The major gas composition of solid body atmospheres is controlled by processes such as volcanic outgassing, surface weathering, and escape to space, whereas much of their trace gas composition is controlled by in situ atmospheric photochemistry. Photochemistry can create hazes that absorb visible and ultraviolet radiation, and it can produce trace gases that might be misconstrued as biosignatures. Methane on Mars and phosphine on Venus, if indeed present, are possible examples. Trace gases can also catalyze reactions that affect major gas composition. For example, the relatively undissociated nature of the upper atmospheres of Mars and Venus is explained by catalytic chemistry that allows CO to recombine with O to reform CO₂ (McElroy and Donahue 1972).

Clouds themselves are of crucial importance in planetary atmospheres. They modify the radiative balance, most significantly in atmospheres with thick, planetwide cloud layers (e.g., sulfuric acid clouds on Venus), although even tenuous clouds can affect the radiative balance of thin atmospheres (e.g., water ice clouds on Mars). Hazes—aerosols formed from photochemical processes in the atmosphere—can also strongly modify radiative balance and may control much of the atmosphere’s vertical thermal structure (Titan and Pluto). In addition, heterogeneous chemistry in clouds and aerosols may be important in some atmospheres (Venus). More data are needed, however, on cloud and haze particle properties and chemical/microphysical formation processes. Figure 9-7 shows clouds and hazes, along with atmospheric temperature profiles, on the four rocky bodies with the densest atmospheres.

Saturn’s largest moon, Titan, is the site of other intriguing atmospheric chemistry. Titan is completely shrouded in organic haze produced from photolysis of methane in its atmosphere. The complicated chemistry leading to its formation is incompletely understood but is thought to include formation of polyacetylenes—long polymers made up of carbon atoms linked by alternating single and triple bonds. Similar chemistry may account for the hazes observed in Pluto’s atmosphere. This chemistry is also of interest to astrobiologists studying the early evolution of life on Earth, as organic haze may have also been present in Earth’s atmosphere during the Archean eon, prior to the rise of atmospheric oxygen.

Q6.6a What Processes Are Important in Controlling the Trace Gas Composition of the Martian Atmosphere?

CH₄ has been reported in Mars’s atmosphere from ground-based spectroscopic measurements (Mumma et al. 2009) and mass spectrometer measurements made by NASA’s Curiosity rover (Webster et al. 2018). These measurements suggest that CH₄ is present at levels of a few parts per billion by volume (ppbv) and that its concentration exhibits both spatial and temporal variations, with maximum values up to ~7 ppbv. But ESA’s Trace Gas Orbiter, which should be capable of detecting CH₄ levels of ~10 parts per trillion, found no evidence for it (Korablev et al. 2019). What is the resolution to this conundrum? Is the methane that Curiosity detects produced locally and destroyed rapidly without having sufficient time to spread throughout the atmosphere? And, if so, is its source biotic or abiotic? A recent report (Civis and Knizek 2021) suggests that CH₄ is produced by ultraviolet-driven surface chemistry, but its rapid removal remains unexplained. An alternative explanation is that methane seeps into Gale Crater from underground sources at night when the planetary boundary layer (i.e., the bottom part of the atmosphere) is stable; then, when daylight arrives, atmospheric mixing dilutes the methane signal below the

detection limits of either the rover or the orbiter (Webster et al. 2021). The mixing ratio of O₂ also shows unexpected seasonal and interannual variability, again suggesting an unknown atmospheric or surface process at work (Trainer et al. 2019). The atmosphere of Mars is expected to contain other trace gases, notably SO₂ and halogens, especially HCl. These gases can provide additional insight into the formation, source, and outgassing mechanisms of methane, as well as any current fumarolic activity on Mars. These gases have not yet been detected on Mars; they could be present at abundances lower than the current detection limit of 0.3 ppbv. Detailed measurements of trace species, especially to determine their isotopic fractions with sufficient accuracy to discriminate between a geological and biological origin, will require the return of a large (in volume or concentration), dedicated sample of atmospheric gas (Beaty et al. 2019).

Q6.6b What Processes Control the Trace Gas Composition of Venus’s Atmosphere at and Below Cloud Level?

Phosphine has been reported in Venus’s atmosphere at cloud-deck level, based on submillimeter-wave observations from Earth (Greaves et al. 2020); however, this detection has been challenged (Snellen et al. 2020). Phosphine is a potential biosignature, so it would be interesting if it were present. Other longstanding questions concerning Venus’s atmospheric composition remain to be answered. What is the mysterious ultraviolet (UV) absorber that blocks roughly half of the solar UV radiation incident on the Venus cloud deck (Esposito 1980)? Sulfur compounds—for example, elemental sulfur chains—are likely candidates, but this hypothesis remains to be tested. Also, what drives the fluctuations of SO₂ in the venusian atmosphere? SO₂ is the main trace gas in the atmosphere and the main progenitor of the sulfuric acid clouds. Observations have detected two-orders-of-magnitude fluctuations of SO₂ in the cloud deck over hours to decades. Do these variations result from volcanic outgassing or from dynamical processes within the atmosphere?

Q6.6c What Photochemical Pathways Take Methane to More Complex Hydrocarbons, Including Hazes, in the Atmospheres of Titan, Triton, Pluto, and Possibly Other Kuiper Belt Objects?

The photolytic destruction of atmospheric methane initiates a chain of photochemical reactions (Hörst 2017) resulting in a plethora of organic species that comprise atmospheric hazes and surface deposits like those observed on Titan by Cassini-Huygens and by ground-based facilities. Similar photochemistry is thought to occur on Triton, Pluto, and perhaps other Kuiper belt objects. This chemistry likely exhibits notable similarities to the chemistry that preceded the emergence of life on Earth. While a great deal about this chemistry has been learned from Cassini-Huygens data, fundamental questions remain.

What is the composition of Titan’s negative ions with mass-to-charge ratios comparable to terrestrial proteins? What are the implications of changes in concentrations of molecules such as HC₃N, C₆H₆, C₄H₂, C₃H₄, and HCN owing to seasonal changes in Titan’s atmosphere (i.e., downwelling in the polar regions of Titan’s stratosphere, leading to enrichment of these species in the stratosphere) and how might this affect the type and extent of different chemical reactions at various levels in the atmosphere? Can abundances of organics detected in Titan’s atmosphere with ground-based telescopes (e.g., C₂H₄, C₄H₂, and C₃H₂) be validated with in situ measurements? What co-crystalline materials or polymorphs might be forming in the haze layers of Titan’s atmosphere as species condense, and how might this inform surface composition and physical properties? What is the composition of the haze in Triton’s atmosphere and how is it expected to evolve with time? And what commonalities exist between the atmospheric hazes at Titan and other bodies such as Pluto (e.g., composition and stratification.)?

Q6.6d How Does Atmospheric Composition, via Ionospheric Chemistry, Affect the Structure and Composition of Exospheres?

Planetary ionospheres and neutral exospheres and thermospheres are closely linked (Schunk and Nagy 2009). Ionization and dissociation, both by solar radiation and by auroral precipitation of energetic charged particles from the external environment, alter the exospheric composition and raise its temperature, on which Jeans escape of

atomic hydrogen and other light species is critically dependent. Ion-neutral chemistry is another key process in the interaction of the neutral atmosphere and exosphere with the ionosphere. At Venus and Mars, ionization of the major neutral species CO₂ produces CO₂⁺ ions, which chemically react with atomic oxygen to produce the main ionospheric species O₂⁺. The O₂⁺ ions dissociatively recombine, generating suprathermal O atoms which populate extensive exospheres at these two planets. At Mars, a large fraction of the atomic oxygen produced then escapes. But the ionospheric chemistry at Venus and Mars has other pathways that are not yet understood and that can produce hot neutral atoms (Bougher et al. 2014). And the nonthermal exospheres of other solar system bodies (e.g., Titan, Triton, Pluto), and their ionospheric source, are poorly understood. The question remains: What superthermal neutral species are produced by ionospheric chemical reactions, thus populating each planetary exosphere?

Q6.6e Do Exospheric Volatile Gases Like Hydrogen, Methane, and Water at the Moon, Mercury, and Other Exposed Solid Bodies Result from Regolith-Grain Chemistry That Converts Atomic Species, Like Solar Wind-Implanted Hydrogen and Carbon, to New Molecules?

In the past decade, there has been a new appreciation that oxygen-rich regolith at exposed rocky bodies can act as chemical conversion surfaces that take in material delivered from the solar wind and micrometeoroids and release new surface-manufactured products into the exosphere. As an example, the Moon has a substantial molecular hydrogen exosphere that results from solar wind proton implantation, hydrogen diffusion, and H₂ formation and release via recombinative desorption (Tucker et al. 2019). The solar wind is collisionless, but the H atoms congregating at the surface of regolith grains interact to form H₂. Methane observed at the Moon by LADEE has also been hypothesized to result from solar wind implantation of ionized hydrogen and carbon (Hodges 2016). LADEE also detected water from the lunar surface during meteor streams, and this has been suggested to be manufactured from impact-related heating of the regolith that is rich in oxygen and implanted (“doped”) hydrogen (Benna et al. 2019). Similar surface chemical processing is expected at Mercury, Phobos, and other solid bodies.

Q6.6f What Processes Control the Formation and Composition of Clouds in the Atmospheres of Venus, Mars, Titan, Triton, and Pluto?

Clouds are seen in the atmospheres of most planetary bodies, from Venus and Mars to Titan, Triton and Pluto (Montmessin et al. 2018), and beyond (see Question 12, Chapter 15). They can affect the atmospheric state via latent heat release, radiative heating, impact on chemistry, and by sequestering atmospheric dust or haze particles as condensation nuclei. However, many uncertainties exist regarding their composition and formation mechanisms. For Venus, better knowledge of cloud processes and composition are needed to explain the mysterious ultraviolet absorbers found there (Q6.6b). On Titan, a major unknown is the distribution, nature, and abundance of the tropospheric clouds over a full annual cycle and whether they are composed of ice, liquid, or a binary methane-nitrogen solution (see also Q6.4e). Studies of the present day coupled dust-volatile cycles of these bodies show that the clouds produced are sensitive to microphysical parameters, such as the heterogeneous nucleation rate of water ice on dust, while the impact on atmospheric state is also sensitive to radiative parameters, controlled at least partly by the size distribution of particles (e.g., Haberle et al. 2019b).

Strategic Research for Q6.6

• Determine the source location and origins of Mars methane by making rapid, accurate measurements of methane fluxes at the surface of Mars on hourly to annual timescales, and by returning an atmospheric sample of sufficient concentration to measure methane isotopic fractions indicative of a biotic or abiotic origin.
• Determine the chemical nature of the unknown ultraviolet absorber in the Venus’s clouds by measuring the composition of the venusian atmosphere and aerosols, particularly at cloud-deck level.

• Constrain the photochemical paths that take methane to more complex hydrocarbons in methane-rich atmospheres (primarily Titan, but also including Triton, Pluto, and other KBOs) by measuring the composition of hazes and aerosols in the atmosphere and on the surface.
• Investigate the microphysical parameters that influence the formation of clouds in planetary atmospheres (primarily Venus, Mars, and Titan, but also including Triton, and Pluto) by determining the distribution, nature and abundance of clouds and the composition and particle size of the droplets comprising them and cloud condensation nuclei around which they form.
• Determine possible pathways for haze formation from experimental study of the chemistry of methane polymerization under both Titan (cold, no oxygen) and early Earth (warm, high CO₂ abundance) conditions.
• Investigate the conversion of regolith-implanted H and C ions into more complex molecules like water, molecular hydrogen, and methane via laboratory studies and molecular dynamic simulations.

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