The Pluto Energetic Particle Spectrometer Science Investigation (PEPSSI) on the New Horizons Mission

The Pluto Energetic Particle Spectrometer Science Investigation (PEPSSI) comprises the hardware and accompanying science investigation on the New Horizons spacecraft to measure pick-up ions from Pluto’s outgassing atmosphere. To the extent that Pluto retains its characteristics similar to those of a “heavy comet” as detected in stellar occultations since the early 1980s, these measurements will characterize the neutral atmosphere of Pluto while providing a consistency check on the atmospheric escape rate at the encounter epoch with that deduced from the atmospheric structure at lower altitudes by the ALICE, REX, and SWAP experiments on New Horizons. In addition, PEPSSI will characterize any extended ionosphere and solar wind interaction while also characterizing the energetic particle environment of Pluto, Charon, and their associated system. First proposed for development for the Pluto Express mission in September 1993, what became the PEPSSI instrument went through a number of development stages to meet the requirements of such an instrument for a mission to Pluto while minimizing the required spacecraft resources. The PEPSSI instrument provides for measurements of ions (with compositional information) and electrons from 10 s of keV to ～1 MeV in a 160°×12° fan-shaped beam in six sectors for 1.5 kg and ～2.5 W.     

Atmospheric Escape and Evolution of Terrestrial Planets and Satellites

The origin and evolution of Venus’, Earth’s, Mars’ and Titan’s atmospheres are discussed from the time when the active young Sun arrived at the Zero-Age-Main-Sequence. We show that the high EUV flux of the young Sun, depending on the thermospheric composition, the amount of IR-coolers and the mass and size of the planet, could have been responsible that hydrostatic equilibrium was not always maintained and hydrodynamic flow and expansion of the upper atmosphere resulting in adiabatic cooling of the exobase temperature could develop. Furthermore, thermal and various nonthermal atmospheric escape processes influenced the evolution and isotope fractionation of the atmospheres and water inventories of the terrestrial planets and Saturn’s large satellite Titan efficiently.     

Aeronomy of Extra-Solar Giant Planets

The intense stellar UV radiation field incident upon extra-solar giant planets causes profound changes to their upper atmospheres. Upper atmospheric temperatures can be tens of thousands of kelvins, causing thermal dissociation of H₂ to H. The stellar ionizing flux converts H to H⁺. The high temperatures also drive large escape rates of H, but for all but the planets with the smallest orbits, this flux is not large enough to affect planet evolution. The escape rate is large enough to drag off heavier atoms such as C and O. For very small orbits, when the hill sphere is inside the atmosphere, escape is unfettered and can affect planet evolution.     

Quantitative Assessments of the Martian Hydrosphere

In this paper, we review current estimates of the global water inventory of Mars, potential loss mechanisms, the thermophysical characteristics of the different reservoirs that water may be currently stored in, and assess how the planet’s hydrosphere and cryosphere evolved with time. First, we summarize the water inventory quantified from geological analyses of surface features related to both liquid water erosion, and ice-related landscapes. They indicate that, throughout most of Martian geologic history (and possibly continuing through to the present day), water was present to substantial depths, with a total inventory ranging from several 100 to as much as 1000 m Global Equivalent Layer (GEL). We then review the most recent estimates of water content based on subsurface detection by orbital and landed instruments, including deep penetrating radars such as SHARAD and MARSIS. We show that the total amount of water measured so far is about 30 m GEL, although a far larger amount of water may be stored below the sounding depths of currently operational instruments. Finally, a global picture of the current state of the subsurface water reservoirs and their evolution is discussed.     

Solar Variability and Climate Impact on Terrestrial Planets

Some possible factors of climate changes and of long term climate evolution are discussed with regard of the three terrestrial planets, Earth, Venus and Mars. Two positive feedback mechanisms involving liquid water, i.e., the albedo mechanism and the greenhouse effect of water vapour, are described. These feedback mechanisms respond to small external forcings, such as resulting from solar or astronomical constants variability, which might thus result in large influences on climatic changes on Earth. On Venus, reactions of the atmosphere with surface minerals play an important role in the climate system, but the involved time scales are much larger. On Mars, climate is changing through variations of the polar axis inclination over time scales of ～10⁵–10⁶ years. Growing evidence also exists that a major climatic change happened on Mars some 3.5 to 3.8 Gigayears ago, leading to the disappearance of liquid water on the planet surface by eliminating most of the CO₂ atmosphere greenhouse power. This change might be due to a large surge of the solar wind, or to atmospheric erosion by large bodies impacts. Indeed, except for their thermospheric temperature response, there is currently little evidence for an effect of long-term solar variability on the climate of Venus and Mars. This fact is possibly due to the absence of liquid water on these terrestrial planets.     

An Overview of Earth’s Global Electric Circuit and Atmospheric Conductivity

The Earth’s global atmospheric electric circuit depends on the upper and lower atmospheric boundaries formed by the ionosphere and the planetary surface. Thunderstorms and electrified rain clouds drive a DC current (～1 kA) around the circuit, with the current carried by molecular cluster ions; lightning phenomena drive the AC global circuit. The Earth’s near-surface conductivity ranges from 10^?7 S?m^?1 (for poorly conducting rocks) to 10^?2 S?m^?1 (for clay or wet limestone), with a mean value of 3.2 S?m^?1 for the ocean. Air conductivity inside a thundercloud, and in fair weather regions, depends on location (especially geomagnetic latitude), aerosol pollution and height, and varies from ～10^?14 S?m^?1 just above the surface to 10^?7 S?m^?1 in the ionosphere at ～80 km altitude. Ionospheric conductivity is a tensor quantity due to the geomagnetic field, and is determined by parameters such as electron density and electron–neutral particle collision frequency. In the current source regions, point discharge (coronal) currents play an important role below electrified clouds; the solar wind-magnetosphere dynamo and the unipolar dynamo due to the terrestrial rotating dipole moment also apply atmospheric potential differences. Detailed measurements made near the Earth’s surface show that Ohm’s law relates the vertical electric field and current density to air conductivity. Stratospheric balloon measurements launched from Antarctica confirm that the downward current density is ～1 pA m^?2 under fair weather conditions. Fortuitously, a Solar Energetic Particle (SEP) event arrived at Earth during one such balloon flight, changing the observed atmospheric conductivity and electric fields markedly. Recent modelling considers lightning discharge effects on the ionosphere’s electric potential (～+250 kV with respect to the Earth’s surface) and hence on the fair weather potential gradient (typically ～130 V?m^?1 close to the Earth’s surface. We conclude that cloud-to-ground (CG) lightning discharges make only a small contribution to the ionospheric potential, and that sprites (namely, upward lightning above energetic thunderstorms) only affect the global circuit in a miniscule way. We also investigate the effects of mesoscale convective systems on the global circuit.     

Venus Atmospheric Thermal Structure and Radiative Balance

From the discovery that Venus has an atmosphere during the 1761 transit by M. Lomonosov to the current exploration of the planet by the Akatsuki orbiter, we continue to learn about the planet’s extreme climate and weather. This chapter attempts to provide a comprehensive but by no means exhaustive review of the results of the atmospheric thermal structure and radiative balance since the earlier works published in Venus and Venus II books from recent spacecraft and Earth based investigations and summarizes the gaps in our current knowledge. There have been no in-situ measurements of the deep Venus atmosphere since the flights of the two VeGa balloons and landers in 1985 (Sagdeev et al., Science 231:1411–1414, 1986). Thus, most of the new information about the atmospheric thermal structure has come from different remote sensing (Earth based and spacecraft) techniques using occultations (solar infrared, stellar ultraviolet and orbiter radio occultations), spectroscopy and microwave, short wave and thermal infrared emissions. The results are restricted to altitudes higher than about 40 km, except for one investigation of the near surface static stability inferred by Meadows and Crisp (J. Geophys. Res. 101:4595–4622, 1996) from 1 \(\upmu\)m observations from Earth. Little information about the lower atmospheric structure is possible below about 40 km altitude from radio occultations due to large bending angles. The gaps in our knowledge include spectral albedo variations over time, vertical variation of the bulk composition of the atmosphere (mean molecular weight), the identity, properties and abundances of absorbers of incident solar radiation in the clouds. The causes of opacity variations in the nightside cloud cover and vertical gradients in the deep atmosphere bulk composition and its impact on static stability are also in need of critical studies. The knowledge gaps and questions about Venus and its atmosphere provide the incentive for obtaining the necessary measurements to understand the planet, which can provide some clues to learn about terrestrial exoplanets.     

Volatile Trapping in Martian Clathrates

Thermodynamic conditions suggest that clathrates might exist on Mars. Despite observations which show that the dominant condensed phases on the surface of Mars are solid carbon dioxide and water ice, clathrates have been repeatedly proposed to play an important role in the distribution and total inventory of the planet’s volatiles. Here we review the potential consequences of the presence of clathrates on Mars. We investigate how clathrates could be a potential source for the claimed existence of atmospheric methane. In this context, plausible clathrate formation processes, either in the close subsurface or at the base of the cryosphere, are reviewed. Mechanisms that would allow for methane release into the atmosphere from an existing clathrate layer are addressed as well. We also discuss the proposed relationship between clathrate formation/dissociation cycles and how potential seasonal variations influence the atmospheric abundances of argon, krypton and xenon. Moreover, we examine several Martian geomorphologic features that could have been generated by the dissociation of extended subsurface clathrate layers. Finally we investigate the future in situ measurements, as well as the theoretical and experimental improvements that will be needed to better understand the influence of clathrates on the evolution of Mars and its atmosphere.     

Long-Term Evolution of the Martian Crust-Mantle System

Lacking plate tectonics and crustal recycling, the long-term evolution of the crust-mantle system of Mars is driven by mantle convection, partial melting, and silicate differentiation. Volcanic landforms such as lava flows, shield volcanoes, volcanic cones, pyroclastic deposits, and dikes are observed on the martian surface, and while activity was widespread during the late Noachian and Hesperian, volcanism became more and more restricted to the Tharsis and Elysium provinces in the Amazonian period. Martian igneous rocks are predominantly basaltic in composition, and remote sensing data, in-situ data, and analysis of the SNC meteorites indicate that magma source regions were located at depths between 80 and 150 km, with degrees of partial melting ranging from 5 to 15 %. Furthermore, magma storage at depth appears to be of limited importance, and secular cooling rates of 30 to 40 K?Gyr^?1 were derived from surface chemistry for the Hesperian and Amazonian periods. These estimates are in general agreement with numerical models of the thermo-chemical evolution of Mars, which predict source region depths of 100 to 200 km, degrees of partial melting between 5 and 20 %, and secular cooling rates of 40 to 50 K?Gyr^?1. In addition, these model predictions largely agree with elastic lithosphere thickness estimates derived from gravity and topography data. Major unknowns related to the evolution of the crust-mantle system are the age of the shergottites, the planet’s initial bulk mantle water content, and its average crustal thickness. Analysis of the SNC meteorites, estimates of the elastic lithosphere thickness, as well as the fact that tidal dissipation takes place in the martian mantle indicate that rheologically significant amounts of water of a few tens of ppm are still present in the interior. However, the exact amount is controversial and estimates range from only a few to more than 200 ppm. Owing to the uncertain formation age of the shergottites it is unclear whether these water contents correspond to the ancient or present mantle. It therefore remains to be investigated whether petrologically significant amounts of water of more than 100 ppm are or have been present in the deep interior. Although models suggest that about 50 % of the incompatible species (H₂O, K, Th, U) have been removed from the mantle, the amount of mantle differentiation remains uncertain because the average crustal thickness is merely constrained to within a factor of two.     

Earthshine Observation of Vegetation and Implication for Life Detection on Other Planets

The detection of exolife is one of the goals of very ambitious future space missions that aim to take direct images of Earth-like planets. While associations of simple molecules present in the planet’s atmosphere (O₂, O₃, CO₂, etc.) have been identified as possible global biomarkers, this paper reviews the detectability of a signature of life from the planet’s surface, i.e. the green vegetation. The vegetation reflectance has indeed a specific spectrum, with a sharp edge around 700 nm, known as the “Vegetation Red Edge” (VRE). Moreover, vegetation covers a large surface of emerged lands, from tropical evergreen forest to shrub tundra. Thus, considering vegetation as a potential global biomarker is relevant. Earthshine allows us to observe the Earth as a distant planet, i.e. without spatial resolution. Since 2001, Earthshine observations have been used by several authors to test and quantify the detectability of the VRE in the Earth spectrum. The vegetation spectral signature is detected as a small “positive shift” of a few percentage points above the continuum, starting at 700 nm. This signature appears in most spectra, and its strength is correlated with the Earth’s phase (visible land versus visible ocean). The observations show that detecting the VRE on Earth requires a photometric relative accuracy of 1% or better. Detecting something equivalent on an Earth-like planet will therefore remain challenging, especially considering the possibility of mineral artifacts and the question of “red edge” universality in the Universe.     

Atmosphere Impact Losses

Determining the origin of volatiles on terrestrial planets and quantifying atmospheric loss during planet formation is crucial for understanding the history and evolution of planetary atmospheres. Using geochemical observations of noble gases and major volatiles we determine what the present day inventory of volatiles tells us about the sources, the accretion process and the early differentiation of the Earth. We further quantify the key volatile loss mechanisms and the atmospheric loss history during Earth’s formation. Volatiles were accreted throughout the Earth’s formation, but Earth’s early accretion history was volatile poor. Although nebular Ne and possible H in the deep mantle might be a fingerprint of this early accretion, most of the mantle does not remember this signature implying that volatile loss occurred during accretion. Present day geochemistry of volatiles shows no evidence of hydrodynamic escape as the isotopic compositions of most volatiles are chondritic. This suggests that atmospheric loss generated by impacts played a major role during Earth’s formation. While many of the volatiles have chondritic isotopic ratios, their relative abundances are certainly not chondritic again suggesting volatile loss tied to impacts. Geochemical evidence of atmospheric loss comes from the \({}^{3}\mathrm{He}/{}^{22}\mathrm{Ne}\), halogen ratios (e.g., F/Cl) and low H/N ratios. In addition, the geochemical ratios indicate that most of the water could have been delivered prior to the Moon forming impact and that the Moon forming impact did not drive off the ocean. Given the importance of impacts in determining the volatile budget of the Earth we examine the contributions to atmospheric loss from both small and large impacts. We find that atmospheric mass loss due to impacts can be characterized into three different regimes: 1) Giant Impacts, that create a strong shock transversing the whole planet and that can lead to atmospheric loss globally. 2) Large enough impactors (\(m_{\mathit{cap}} \gtrsim \sqrt{2} \rho_{0} (\pi h R)^{3/2}\), \(r_{\mathit{cap}}\sim25~\mbox{km}\) for the current Earth), that are able to eject all the atmosphere above the tangent plane of the impact site, where \(h\), \(R\) and \(\rho_{0}\) are the atmospheric scale height, radius of the target, and its atmospheric density at the ground. 3) Small impactors (\(m_{\mathit{min}}>4 \pi\rho_{0} h^{3}\), \(r_{\mathit {min}}\sim 1~\mbox{km}\) for the current Earth), that are only able to eject a fraction of the atmospheric mass above the tangent plane. We demonstrate that per unit impactor mass, small impactors with \(r_{\mathit{min}} < r < r_{\mathit{cap}}\) are the most efficient impactors in eroding the atmosphere. In fact for the current atmospheric mass of the Earth, they are more than five orders of magnitude more efficient (per unit impactor mass) than giant impacts, implying that atmospheric mass loss must have been common. The enormous atmospheric mass loss efficiency of small impactors is due to the fact that most of their impact energy and momentum is directly available for local mass loss, where as in the giant impact regime a lot of energy and momentum is ’wasted’ by having to create a strong shock that can transverse the entirety of the planet such that global atmospheric loss can be achieved. In the absence of any volatile delivery and outgassing, we show that the population of late impactors inferred from the lunar cratering record containing 0.1% \(M_{\oplus }\) is able to erode the entire current Earth’s atmosphere implying that an interplay of erosion, outgassing and volatile delivery is likely responsible for determining the atmospheric mass and composition of the early Earth. Combining geochemical observations with impact models suggest an interesting synergy between small and big impacts, where giant impacts create large magma oceans and small and larger impacts drive the atmospheric loss.     

A Comparative Study of the Influence of the Active Young Sun on the Early Atmospheres of Earth, Venus, and Mars

Because the solar radiation and particle environment plays a major role in all atmospheric processes such as ionization, dissociation, heating of the upper atmospheres, and thermal and non-thermal atmospheric loss processes, the long-time evolution of planetary atmospheres and their water inventories can only be understood within the context of the evolving Sun. We compare the effect of solar induced X-ray and EUV (XUV) heating on the upper atmospheres of Earth, Venus and Mars since the time when the Sun arrived at the Zero-Age-Main-Sequence (ZAMS) about 4.6 Gyr ago. We apply a diffusive-gravitational equilibrium and thermal balance model for studying heating of the early thermospheres by photodissociation and ionization processes, due to exothermic chemical reactions and cooling by IR-radiating molecules like CO₂, NO, OH, etc. Our model simulations result in extended thermospheres for early Earth, Venus and Mars. The exospheric temperatures obtained for all the three planets during this time period lead to diffusion-limited hydrodynamic escape of atomic hydrogen and high Jeans’ escape rates for heavier species like H₂, He, C, N, O, etc. The duration of this blow-off phase for atomic hydrogen depends essentially on the mixing ratios of CO₂, N₂ and H₂O in the atmospheres and could last from ∼100 to several hundred million years. Furthermore, we study the efficiency of various non-thermal atmospheric loss processes on Venus and Mars and investigate the possible protecting effect of the early martian magnetosphere against solar wind induced ion pick up erosion. We find that the early martian magnetic field could decrease the ion-related non-thermal escape rates by a great amount. It is possible that non-magnetized early Mars could have lost its whole atmosphere due to the combined effect of its extended upper atmosphere and a dense solar wind plasma flow of the young Sun during about 200 Myr after the Sun arrived at the ZAMS. Depending on the solar wind parameters, our model simulations for early Venus show that ion pick up by strong solar wind from a non-magnetized planet could erode up to an equivalent amount of ∼250 bar of O⁺ ions during the first several hundred million years. This accumulated loss corresponds to an equivalent mass of ∼1 terrestrial ocean (TO (1 TO ∼1.39×10²⁴ g or expressed as partial pressure, about 265 bar, which corresponds to ∼2900 m average depth)). Finally, we discuss and compare our findings with the results of preceding studies.     

Water and Brines on Mars: Current Evidence and Implications for MSL

Liquid water is a basic ingredient for life as we know it. Therefore, in order to understand the habitability of other planets we must first understand the behavior of water on them. Mars is the most Earth-like planet in the solar system and it has large reservoirs of H₂O. Here, we review the current evidence for pure liquid water and brines on Mars, and discuss their implications for future and current missions such as the Mars Science Laboratory. Neither liquid water nor liquid brines are currently stable on the surface of Mars, but they could be present temporarily in a few areas of the planet. Pure liquid water is unlikely to be present, even temporarily, on the surface of Mars because evaporation into the extremely dry atmosphere would inhibit the formation of the liquid phase, where the temperature and pressure are high enough so that water would neither freeze nor boil. The exception to this is that monolayers of liquid water, referred to as undercooled liquid interfacial water, could exist on most of the Martian surface. In a few places liquid brines could exist temporarily on the surface because they could form at cryogenic temperatures, near ice or frost deposits where sublimation could be inhibited by the presence of nearly saturated air. Both liquid water and liquid brines might exist in the shallow subsurface because even a thin layer of soil forms an effective barrier against sublimation allowing pure liquid water to form sporadically in a few places, or liquid brines to form over longer periods of time in large portions of the planet. At greater depths, ice deposits could melt where the soil conductivity is low enough to blanket the deeper subsurface effectively. This could cause the formation of aquifers if the deeper soil is sufficiently permeable and an impermeable layer exists below the source of water. The fact that liquid brines and groundwater are likely to exist on Mars has important implications for geochemistry, glaciology, mineralogy, weathering and the habitability of Mars.     

Planetary Magnetic Fields and Solar Forcing: Implications for Atmospheric Evolution

The solar wind and the solar XUV/EUV radiation constitute a permanent forcing of the upper atmosphere of the planets in our solar system, thereby affecting the habitability and chances for life to emerge on a planet. The forcing is essentially inversely proportional to the square of the distance to the Sun and, therefore, is most important for the innermost planets in our solar system—the Earth-like planets. The effect of these two forcing terms is to ionize, heat, chemically modify, and slowly erode the upper atmosphere throughout the lifetime of a planet. The closer to the Sun, the more efficient are these process. Atmospheric erosion is due to thermal and non-thermal escape. Gravity constitutes the major protection mechanism for thermal escape, while the non-thermal escape caused by the ionizing X-rays and EUV radiation and the solar wind require other means of protection. Ionospheric plasma energization and ion pickup represent two categories of non-thermal escape processes that may bring matter up to high velocities, well beyond escape velocity. These energization processes have now been studied by a number of plasma instruments orbiting Earth, Mars, and Venus for decades. Plasma measurement results therefore constitute the most useful empirical data basis for the subject under discussion. This does not imply that ionospheric plasma energization and ion pickup are the main processes for the atmospheric escape, but they remain processes that can be most easily tested against empirical data. Shielding the upper atmosphere of a planet against solar XUV, EUV, and solar wind forcing requires strong gravity and a strong intrinsic dipole magnetic field. For instance, the strong dipole magnetic field of the Earth provides a “magnetic umbrella”, fending of the solar wind at a distance of 10 Earth radii. Conversely, the lack of a strong intrinsic magnetic field at Mars and Venus means that the solar wind has more direct access to their topside atmosphere, the reason that Mars and Venus, planets lacking strong intrinsic magnetic fields, have so much less water than the Earth? Climatologic and atmospheric loss process over evolutionary timescales of planetary atmospheres can only be understood if one considers the fact that the radiation and plasma environment of the Sun has changed substantially with time. Standard stellar evolutionary models indicate that the Sun after its arrival at the Zero-Age Main Sequence (ZAMS) 4.5 Gyr ago had a total luminosity of ≈70% of the present Sun. This should have led to a much cooler Earth in the past, while geological and fossil evidence indicate otherwise. In addition, observations by various satellites and studies of solar proxies (Sun-like stars with different age) indicate that the young Sun was rotating more than 10 times its present rate and had correspondingly strong dynamo-driven high-energy emissions which resulted in strong X-ray and extreme ultraviolet (XUV) emissions, up to several 100 times stronger than the present Sun. Further, evidence of a much denser early solar wind and the mass loss rate of the young Sun can be determined from collision of ionized stellar winds of the solar proxies, with the partially ionized gas in the interstellar medium. Empirical correlations of stellar mass loss rates with X-ray surface flux values allows one to estimate the solar wind mass flux at earlier times, when the solar wind may have been more than 1000 times more massive. The main conclusions drawn on basis of the Sun-in-time-, and a time-dependent model of plasma energization/escape is that:

1. Solar forcing is effective in removing volatiles, primarily water, from planets,
2. planets orbiting close to the early Sun were subject to a heavy loss of water, the effect being most profound for Venus and Mars, and
3. a persistent planetary magnetic field, like the Earth’s dipole field, provides a shield against solar wind scavenging.

     

Planetary Magnetic Dynamo Effect on Atmospheric Protection of Early Earth and Mars

In light of assessing the habitability of Mars, we examine the impact of the magnetic field on the atmosphere. When there is a magnetic field, the atmosphere is protected from erosion by solar wind. The magnetic field ensures the maintenance of a dense atmosphere, necessary for liquid water to exist on the surface of Mars. We also examine the impact of the rotation of Mars on the magnetic field. When the magnetic field of Mars ceased to exist (about 4 Gyr ago), atmospheric escape induced by solar wind began. We consider scenarios which could ultimately lead to a decrease of atmospheric pressure to the presently observed value of 7 mbar: a much weaker early martian magnetic field, a late onset of the dynamo, and high erosion rates of a denser early atmosphere.     

Composition and Measurement of Charged Atmospheric Clusters

Atmospheric charged clusters are formed in a series of rapid chemical reactions after ionisation, leaving a central ion X⁺ or X^? clustered with n ligands (Y) _n . In solar system tropospheres and stratospheres there are two distinct cluster regimes: the terrestrial planets contain largely hydrated clusters (i.e. Y=H₂O), whereas the gas planets and their moons have organic or nitrogenated cluster species. These classifications are largely based on model predictions, since hardly any measurements are available. The few existing composition measurements are reviewed, including the recent detection of massive charged particles in Titan’s upper atmosphere. Technologies for both remote sensing and in situ measurements of atmospheric charged clusters are discussed. Preliminary measurements in the terrestrial atmosphere are presented indicating that ambient charged cluster species interact with downwelling infra-red radiation at 9.15 μm, even in the presence of cloud. This supports the possibility of future infrared detection of charged clusters.     

Infrared Spectroscopy of Solar-System Planets

Most of our knowledge regarding planetary atmospheric composition and structure has been achieved by remote sensing spectroscopy. Planetary spectra strongly differ from one planet to another. CO₂ signatures dominate on Mars, and even more on Venus (where the thermal component is detectable down to 1 μm on the dark side). Spectroscopic monitoring of Venus, Earth and Mars allows us to map temperature fields, wind fields, clouds, aerosols, surface mineralogy (in the case of the Earth and Mars), and to study the planets’ seasonal cycles. Spectra of giant planets are dominated by H₂, CH₄ and other hydrocarbons, NH₃, PH₃ and traces of other minor compounds like CO, H₂O and CO₂. Measurements of the atmospheric composition of giant planets have been used to constrain their formation scenario.     

Atmospheric/Exospheric Characteristics of Icy Satellites

The atmospheres/exospheres of icy satellites greatly vary from one to the next in terms of density, composition, structure or steadiness. Titan is the only icy satellite with a dense atmosphere comparable in many ways to that of the Earth’s atmosphere. Titan’s atmosphere prevents the surface from direct interaction with the plasma environment, but gives rise to Earth-like exchanges of energy, matter and momentum. The atmospheres of other satellites are tenuous. Enceladus’ atmosphere manifests itself in a large water vapor plume emanating from surface cracks near the south pole. Io’s SO₂ atmosphere originates from volcanoes. Europa’s tenuous O₂ atmosphere is produced by intense radiation bombardment. This chapter reviews the characteristics of the atmospheres of Titan, Enceladus, Io and Europa based on observations.     

Geochemical Reservoirs and Timing of Sulfur Cycling on Mars

Sulfate-dominated sedimentary deposits are widespread on the surface of Mars, which contrasts with the rarity of carbonate deposits, and indicates surface waters with chemical features drastically different from those on Earth. While the Earth’s surface chemistry and climate are intimately tied to the carbon cycle, it is the sulfur cycle that most strongly influences the Martian geosystems. The presence of sulfate minerals observed from orbit and in-situ via surface exploration within sedimentary rocks and unconsolidated regolith traces a history of post-Noachian aqueous processes mediated by sulfur. These materials likely formed in water-limited aqueous conditions compared to environments indicated by clay minerals and localized carbonates that formed in surface and subsurface settings on early Mars. Constraining the timing of sulfur delivery to the Martian exosphere, as well as volcanogenic H₂O is therefore central, as it combines with volcanogenic sulfur to produce acidic fluids and ice. Here, we reassess and review the Martian geochemical reservoirs of sulfur from the innermost core, to the mantle, crust, and surficial sediments. The recognized occurrences and the mineralogical features of sedimentary sulfate deposits are synthesized and summarized. Existing models of formation of sedimentary sulfate are discussed and related to weathering processes and chemical conditions of surface waters. We also review existing models of sulfur content in the Martian mantle and analyze how volcanic activities may have transferred igneous sulfur into the exosphere and evaluate the mass transfers and speciation relationships between volcanic sulfur and sedimentary sulfates. The sedimentary clay-sulfate succession can be reconciled with a continuous volcanic eruption rate throughout the Noachian-Hesperian, but a process occurring around the mid-Noachian must have profoundly changed the composition of volcanic degassing. A hypothetical increase in the oxidation state or in water content of Martian lavas or a decrease in atmospheric pressure is necessary to account for such a change in composition of volcanic gases. This would allow the pre mid-Noachian volcanic gases to be dominated by water and carbon-species but late Noachian and Hesperian volcanic gases to be sulfur-rich and characterized by high SO₂ content. Interruption of early dynamo and impact ejection of the atmosphere may have decreased the atmospheric pressure during the early Noachian whereas it remains unclear how the redox state or water content of lavas could have changed. Nevertheless, volcanic emission of SO₂ rich gases since the late Noachian can explain many features of Martian sulfate-rich regolith, including the mass of sulfate and the particular chemical features (i.e. acidity) of surface waters accompanying these deposits. How SO₂ impacted on Mars’s climate, with possible short time scale global warming and long time scale cooling effects, remains controversial. However, the ancient wet and warm era on Mars seems incompatible with elevated atmospheric sulfur dioxide because conditions favorable to volcanic SO₂ degassing were most likely not in place at this time.