Venus is a slowly rotating planet with a thick atmosphere (~9.2 MPa at the surface). Ground- and satellite-based observations have shown atmospheric superrotation (atmospheric rotation much faster than solid surface rotation), global-scale cloud patterns (e.g., Y-shaped and bow-shaped structures), and polar vortices (polar hot dipole and fine structures). The Venusian atmospheric circulation, controlled by the planet’s radiative forcing and astronomical parameters, is quite different from the earth’s. As the meteorological data have been stored, understanding of the atmospheric circulation has been gradually enriched with the help of theories of geophysical fluid dynamics and meteorology.
In the cloud layer far from the surface (49–70 km altitude), superrotational flows (east-to-west zonal winds) exceeding 100 m/s and meridional (equator-to-pole) flows have been observed along with planetary-scale brightness variations unique to Venus. The fully developed superrotation, which is ~60 times faster than the planetary rotation, is maintained by meridional circulation and waves. For the planetary-scale variations, slow-traveling waves with stationary and solar-locked structures and fast-traveling waves with phase velocities of around the superroational wind speeds are dominant in the cloud layer. Thermal tides, Rossby waves, Kelvin waves, and gravity waves play important roles in mechanisms for maintaining fast atmospheric rotation. In the lower atmosphere below the cloud layer, the atmospheric circulation is still unknown because of the lack of global observations. In addition to the limited observations, the atmospheric modeling contributes to deep understanding of the atmospheric circulation system. Recent general circulation models have well simulated the dynamical and thermal structures of Venus’s atmosphere, though there remain outstanding issues.
Karen Aplin and Georg Fischer
Electricity occurs in atmospheres across the Solar System planets and beyond, spanning spectacular lightning displays in clouds of water or dust, to more subtle effects of charge and electric fields. On Earth, lightning is likely to have existed for a long time, on the basis of evidence from fossilized lightning strikes in ancient rocks, but observations of planetary lightning are necessarily much more recent. The generation and observations of lightning and other atmospheric electrical processes, both from within-atmosphere measurements, and spacecraft remote sensing, can be readily studied using a comparative planetology approach, with the Earth as a model.
All atmospheres contain charged molecules, electrons, and/or molecular clusters created by ionization from cosmic rays and other processes, which may affect an atmosphere’s energy balance both through aerosol and cloud formation and direct absorption of radiation. Several planets are anticipated to host a “global electric circuit” by analogy with the circuit occurring on the Earth, where thunderstorms drive the current of ions or electrons through weakly conductive parts of the atmosphere. This current flow may further modulate an atmosphere’s radiative properties through cloud and aerosol effects.
Lightning could potentially have implications for life through its effects on atmospheric chemistry and particle transport. It has been observed on many of the Solar System planets (Earth, Jupiter, Saturn, Uranus, and Neptune), and it may also be present on Venus and Mars. On Earth, Jupiter, and Saturn, lightning is thought to be generated in deep water and ice clouds, but discharges can be generated in dust, as for terrestrial volcanic lightning, and on Mars. Other, less well-understood mechanisms causing discharges in non-water clouds also seem likely. The discovery of thousands of exoplanets has recently led to a range of further exotic possibilities for atmospheric electricity, though lightning detection beyond our Solar System remains a technical challenge to be solved.
Chemical and phase compositions of the surface of Venus could reflect a history of gas–rock and fluid–rock interactions, recent and past climate changes, and a loss of water from the Earth’s sister planet. The concept of chemical weathering on Venus through gas–solid type reactions was established in the early 1960s after the discovery of the hot and dense CO2-rich atmosphere of the planet, inferred from Earth-based and Mariner 2 radio emission data. Initial models suggested carbonation, hydration, and oxidation of exposed igneous rocks and a control (buffering) of atmospheric gases by solid–gas type chemical equilibria in the near-surface rocks. Carbonates, phyllosilicates and Fe oxides were considered likely secondary minerals. From the late 1970s onward, measurements of trace gases in the sub-cloud atmosphere by the Pioneer Venus and Venera entry probes and by Earth-based infrared spectroscopy challenged the likelihood of hydration and carbonation. The atmospheric H2O gas content appeared to be low enough to allow the stable existence of H2O-bearing and a majority of OH-bearing minerals. The concentration of SO2 gas was too high to allow the stability of Ca-rich carbonates and silicates with respect to sulfatization to CaSO4. In the 1980s, the detection of an elevated bulk S content at the Venera and Vega landing sites suggested ongoing consumption of atmospheric SO2 to surface sulfates. The supposed composition of the near-surface atmosphere implied oxidation of ferrous minerals to Fe oxides, magnetite and hematite, consistent with the infrared reflectance of surface materials. The likelihood of sulfatization and oxidation has been illustrated in modeling experiments in simulated Venus’ conditions. The morphology of Venus’ surface suggests contact of atmospheric gases with hot surface materials of mainly basaltic composition during the several hundreds of millions years since a global volcanic/tectonic resurfacing. Some exposed materials could have reacted at higher and lower temperatures in a presence of diverse gases at different altitudinal, volcanic, impact, and atmospheric settings. On highly deformed tessera terrains, more ancient rocks of unknown composition may reflect interactions with putative water-rich atmospheres and even aqueous solutions. Geological formations rich in salt, carbonate, Fe oxide, or silica will indicate past aqueous processes. The apparent diversity of affected solids, surface temperatures, pressures, and gas/fluid compositions throughout Venus’ history implies multiple signs of chemical alterations that remain to be investigated. The current understanding of chemical weathering is limited by the uncertain composition of the deep atmosphere, by the lack of direct data on the phase and chemical composition of surface materials, and by the uncertain data on thermodynamics of minerals and their solid solutions. In preparation for further atmospheric entry probe and lander missions, rock alteration could be investigated through chemical kinetic experiments and calculations of solid-gas/fluid equilibria to constrain past and present processes.
Anni Määttänen and Franck Montmessin
This is an advance summary of a forthcoming article in the Oxford Research Encyclopedia of Planetary Science. Please check back later for the full article.
Although resembling an extremely dry desert, planet Mars hosts clouds in its atmosphere. Every day somewhere on the planet a part of the tiny amount of water vapor held by the atmosphere can condense as ice crystals to form cirrus-type clouds. The existence of water ice clouds has been known for a long time, and they have been studied for decades, leading to the establishment of a well-known climatology and understanding of their formation and properties. Despite their thinness, they have a clear impact on the atmospheric temperatures, thus affecting the Martian climate.
Another, more exotic type of clouds forms as well on Mars. The atmospheric temperatures can plunge to such frigid values that the major gaseous component of the atmosphere, CO2, condenses as ice crystals. These clouds form in the cold polar night where they also contribute to the formation of the CO2 ice polar cap, and also in the mesosphere at very high altitudes, near the edge of space, analogously to the noctilucent clouds on Earth. The mesospheric clouds are a fairly recent discovery and have put our understanding of the Martian atmosphere to a test.
On Mars, cloud crystals form on ice nuclei, mostly provided by the omnipresent dust. Thus, the clouds link the three major climatic cycles: those of the two major volatiles, H2O and CO2; and that of dust, which is a major climatic agent itself.
Matthew R. Balme
Dust devils are rotating columns or cones of air, loaded with dust and other fine particles, that are most often found in arid or desert areas. They are common on both Mars and Earth, despite Mars’ very thin atmosphere. The smallest and least intense dust devils might last only a few 10s of seconds and be just a meters or two across. The largest dust devils can persist for hours and are intensely swirling columns of dust with “skirts” of sand at their base, 10s or more meters in diameter and hundreds of meters high; even larger examples have been seen on Mars. Dust devils on Earth have been documented for thousands of years, but scientific observations really began in the early 20th century, culminating in a period of intense research in the 1960s. The discovery of dust devils on Mars was made using data from the NASA Viking lander and orbiter missions in the late 1970s and early 1980s and stimulated a renewed scientific interest in dust devils. Observations from subsequent lander, rover, and orbital missions show that Martian dust devils are common but heterogeneously distributed in space and time and have a significant effect on surface albedo (often leaving “tracks” on the surface) but do not appear to be triggers of global or major dust storms. An aspiration of future research is to synthesize observations and detailed models of dust devils to estimate more accurately their role in dust lifting at both local and global scales, both on Earth and on Mars.
Dmitry V. Bisikalo, Pavel V. Kaygorodov, and Valery I. Shematovich
The history of exoplanetary atmospheres studies is strongly based on the observations and investigations of the gaseous envelopes of hot Jupiters—exoplanet gas giants that have masses comparable to the mass of Jupiter and orbital semi-major axes shorter than 0.1 AU. The first exoplanet around a solar-type star was a hot Jupiter discovered in 1995. Researchers found an object that had completely atypical parameters compared to planets known in the solar system. According to their estimates, the object might have a mass about a half of the Jovian mass and a very short orbital period (four days), which means that it has an orbit roughly corresponding to the orbit of Mercury. Later, many similar objects were discovered near different stars, and they acquired a common name—hot Jupiters. It is still unclear what the mechanism is for their origin, because generally accepted theories of planetary evolution predict the formation of giant planets only at large orbital distances, where they can accrete enough matter before the protoplanetary disc disappears. If this is true, before arriving at such low orbits, hot Jupiters might have a long migration path, caused by interactions with other massive planets and/or with the gaseous disc. In favor of this model is the discovery of many hot Jupiters in elliptical and highly inclined orbits, but on the other hand several observed hot Jupiters have circular orbits with low inclination. An alternative hypothesis is that the cores of future hot Jupiters are super-Earths that may later intercept matter from the protoplanetary disk falling on the star.
The scientific interest in hot Jupiters has two aspects. The first is the peculiarity of these objects: they have no analogues in the solar system. The second is that, until recently, only for hot Jupiters was it possible to obtain observational characteristics of their atmospheres. Many of the known hot Jupiters are eclipsing their host stars, so, from their light curve and spectral data obtained during an eclipse, it became possible to obtain information about their shape and their atmospheric composition.
Thus it is possible to conclude that hot Jupiters are a common type of exoplanet, having no analogues in the solar system. Many aspects of their evolution and internal structure remain unclear. Being very close to their host stars, hot Jupiters must interact with the stellar wind and stellar magnetic field, as well as with stellar flares and coronal mass ejections, allowing researchers to gather information about them. According to UV observations, at least a fraction of hot Jupiters have extended gaseous envelopes, extending far beyond of their upper atmospheres. The envelopes are observable with current astronomical instruments, so it is possible to develop their astrophysical models. The history of hot Jupiter atmosphere studies during the past 20 years and the current status of modern theories describing the extended envelopes of hot Jupiters are excellent examples of the progress in understanding planetary atmospheres formation and evolution both in the solar system and in the extrasolar planetary systems.
Valery I. Shematovich and Dmitry V. Bisikalo
This is an advance summary of a forthcoming article in the Oxford Research Encyclopedia of Planetary Science. Please check back later for the full article.
The uppermost layers of a planetary atmosphere, where the density of neutral particles is vanishingly low, are commonly called the exosphere or the planetary corona. Since the atmosphere is not completely bound to the planet by the planetary gravitational field, light atoms, such as hydrogen and helium with sufficiently large velocities, can escape from the upper atmosphere into interplanetary space. This process is commonly called Jeans escape, and it depends on the temperature of the ambient atmospheric gas at an altitude where the atmospheric gas is virtually collisionless. The heavier carbon, nitrogen, and oxygen atoms can escape from the atmospheres of the terrestrial planets only through non-thermal processes such as photo- and electron-impact dissociation, charge exchange, atmospheric sputtering, and ion pick-up. Theories of planetary exospheres have been based on ground-based and space observations of emission features such as the 121.6 nm Ly-α and 102.6 nm Ly-β hydrogen lines, the 58.4 nm helium line, and the 130.4 and 135.6 nm atomic oxygen lines. Such observations, together with in situ mass-spectrometer measurements, as at Titan, allow the density and temperature height profiles of the exospheric components to be constructed. The measurements reveal that planetary coronas contain both a fraction of thermal neutral particles with a mean kinetic energy corresponding to the exospheric temperature and a fraction of hot neutral particles with mean kinetic energy much higher than the exospheric temperature. These suprathermal (hot) atoms and molecules are a direct manifestation of the non-thermal processes taking place in the atmospheres. These hot particles lead to the atmospheric escape, determine the coronal structure, produce non-thermal emissions, and react with the ambient atmospheric gas triggering hot atom chemistry.
One of the brightest manifestations of these processes is a formation of hot oxygen corona around terrestrial planets. Oxygen atom is one of the lightest among heavy atmospheric species, so it is a best species to form corona, and another important aspect is that it produces a lot of observational evidence. The transport of suprathermal oxygen atoms to exospheric heights leads to the formation of hot oxygen coronas around Venus, Earth, and Mars. It has been well established by both observations and theoretical calculations that hot oxygen is an important constituent in the transition region between upper thermosphere and exosphere at terrestrial planets.
The study of the planetary coronas is based on direct observations and numerical simulations. It is a rarefied gas, therefore, production and transport of suprathermal particles into the corona requires solving a Boltzmann equation or a DSMC simulation. The stochastic simulation method had been widely used to investigate the formation, kinetics, and transport of suprathermal particles in the hot planetary coronas. This approach was first used to study the formation of the hot oxygen geocorona, taking into account the exothermic chemistry and the precipitation of magnetospheric protons and high-energy O+ ions from the ring current. It was found that only atmospheric sputtering results in the formation of the escape flux of energetic oxygen atoms, providing an important source of heavy atoms for the magnetosphere and geospace. A stochastic modeling approach was also applied to study the escape of hot oxygen atoms from the upper atmosphere of Mars and Venus; the kinetics and transport of suprathermal atoms and molecules in the hot oxygen corona at Jovian satellite Europa, which is an example of a highly non-equilibrium near-surface atmosphere; and the hot extended corona at Saturnian satellite Titan, which was directly measured by the spacecraft Cassini.
Timothy E. Dowling
Jet streams, “jets” for short, are remarkably coherent streams of air found in every major atmosphere. They have a profound effect on a planet’s global circulation and have been an enigma since the belts and zones of Jupiter were discovered in the 1600s. Collaborations between observers, experimentalists, computer modelers, and applied mathematicians seek to understand what processes affect jet size, strength, direction, shear stability, and predictability. Key challenges include nonlinearity, nonintuitive wave physics, nonconstant-coefficient differential equations, and the many nondimensional numbers that arise from the competing physical processes that affect jets, including gravity, pressure gradients, Coriolis accelerations, and turbulence. Fortunately, the solar system provides many examples of jets, and both laboratory and computer simulations allow for carefully controlled experiments. Jet research is multidisciplinary but is united by a common language, the conservation of potential vorticity (PV), which is an all-in-one conservation law that combines the conservation laws of mass, momentum, and thermal energy into a single expression. The leading theories of how jets emerge out of turbulence, and why they are invariably zonal (east-west orientated), reveal the importance of vorticity waves that owe their existence to conservation of PV.
Jets are observed to naturally group into equatorial, midlatitude, and polar types. Earth and Uranus have weakly retrograde equatorial jets, but most planets exhibit strongly prograde (superrotating) equatorial jets, which require eddies to transport momentum up-gradient in a manner that is not obvious but is beginning to be understood. Jupiter and Saturn exhibit multiple alternating jets spanning their midlatitudes, with deep roots that connect to their interior circulations. Polar jets universally exhibit an impressive inhibition of meridional (north-south) mixing, and the seasonal nature of the polar jets on Earth, Mars, and Titan contrasts with the permanence of those on the giant planets, including Saturn’s beautiful north-polar hexagon. Intriguingly, jets in atmospheres have strong analogies with jets in nonneutral plasmas, with practical benefits to both disciplines.
The Martian ionosphere is a plasma embedded within the neutral upper atmosphere of the planet. Its main source is the ionization of the CO2-dominated Martian mesosphere and thermosphere by energetic EUV solar radiation. The ionosphere of Mars is subject to an important variability induced by changes in its forcing mechanisms (e.g., the UV solar flux) and by variations in the neutral atmosphere (e.g., the presence of global dust storms, atmospheric waves and tides, changes in atmospheric composition, etc.). Its vertical structure is dominated by a maximum in electron concentration at altitude about 120–140 km, coincident with the peak of the ionization rate. Below, there is a secondary peak produced by solar X-rays and photoelectron-impact ionization. A sporadic third layer, possibly of meteoric origin, has been also detected below.
The most abundant ion in the Martian ionosphere is O2 +, although O+ can become more abundant in the upper ionospheric layers. While below about 180–200 km the Martian ionosphere is dominated by photochemical processes, above those altitudes the dynamics of the plasma becomes more important. The ionosphere is also an important source of escaping particles via processes such as dissociative recombination of ions or ion pickup. So, characterization of the ionosphere provides or can provide information about such disparate systems and processes as solar radiation reaching the planet, the neutral atmosphere, meteoric influx, atmospheric escape to space, or the interaction of the planet with the solar wind.
It is thus not surprising that the interest about this region dates from the beginning of the space era. From the first measurements provided by the Mariner 4 mission in the 1960s to observations by the Mars Express and MAVEN orbiters in the 2010s, our knowledge of this atmospheric region is the consequence of the accumulation of more than 50 years of discontinuous measurements by different space missions. Numerical simulations by computational models able to simulate the processes that shape the ionosphere have also been commonly employed to obtain information about this region, to provide an interpretation of the observations and to fill their gaps. As a result, at the end of the 2010s the Martian ionosphere was the best known one after that of the Earth. However, there are still areas for which our knowledge is far from being complete. Examples are the details and balance of the mechanisms populating the nightside ionosphere, the origin and variability of the lower ionospheric peak, and the precise mechanisms shaping the topside ionosphere.
The planetary boundary layer of Mars is a crucial component of the Martian climate and meteorology, as well as a key driver of the surface-atmosphere exchanges on Mars. As such, it is explored by several landers and orbiters; high-resolution atmospheric modeling is used to interpret the measurements by those spacecrafts. The planetary boundary layer of Mars is particularly influenced by the strong radiative control of the Martian surface and, as a result, features a more extreme version of planetary boundary layer phenomena occurring on Earth. In daytime, the Martian planetary boundary layer is highly turbulent, mixing heat and momentum in the atmosphere up to about 10 kilometers from the surface. Daytime convective turbulence is organized as convective cells and vortices, the latter giving rise to numerous dust devils when dust is lifted and transported in the vortex. The nighttime planetary boundary layer is dominated by stable-layer turbulence, which is much less intense than in the daytime, and slope winds in regions characterized by uneven topography. Clouds and fogs are associated with the planetary boundary layer activity on Mars.
Angela M. Zalucha and Jason Cook
In addition to ground-based observations beginning in the 1970s, NASA’s Voyager 2 spacecraft flew by Triton in 1989, and NASA’s New Horizons spacecraft flew by Pluto in 2015. Prior to the flyby of New Horizons, Pluto and Triton were termed “sister worlds” due to what appeared to be a high degree of similarity in solid-body density, surface ices, diameter, and surface pressures. Despite being small, cold, icy bodies, both Pluto and Triton have been found to have atmospheres that behave as a continuous fluid up to 300 km altitude above the surface and thereby have a defined temperature, surface pressure, and global general circulation (wind). The primary constituent of these atmospheres is molecular nitrogen, with methane and carbon monoxide comprising the largest abundances of trace gases. The surface pressure as measured in the 2010s on both worlds is of the order of 10 microbars (1 Pa = 10 µbar), for these exotic atmospheres exchange mass between sublimation of surface ice and deposition of nitrogen over the course of each body’s year. Ground-based stellar occultation measurements observed a dramatic change in surface pressure, which one study found was as much as a factor of two increase between 1988 and 2003 on Pluto, presumably due to Pluto’s seasonal volatile cycle. Voyager 2 observed plumes and surface “streaks” on Triton, while New Horizons observed dunes (indicating wind speeds of 1–10 m s−1) as well as streaks, evidently indicating the presence of surface and near-surface winds.
While wind velocity aloft has not been directly measured on Pluto or Triton, 3-D general circulation modeling studies of both worlds have shown zonal (east–west) wind speeds of the order of 10 m/s, meridional (north–south) wind speeds of the order of 1 m/s, and extremely weak vertical wind speeds.
In 2015, New Horizons showed that Pluto and Triton were much more different than previously thought. New Horizons uncovered many spectacular views of Pluto’s atmosphere. First, while hydrocarbon haze was observed on Triton, Pluto had multiple, very distinct stratified haze layers bearing a similar appearance to the layers of an onion. Second, Pluto’s surface elevation was found to be largely inhomogeneous (in contrast to Triton) in the form of a large depression (Sputnik Planitia). Third, the characteristics of the surface markings on Pluto were found to be different than the streaks observed on Triton, which has implications for surface wind patterns.
Further major discoveries made by New Horizons included evidence for many hydrocarbon species in trace concentrations, a lower than expected surface pressure, which could previously only be indirectly ascertained from ground-based observations, and a higher mixing ratio of methane at higher altitudes than at lower due to gravitational diffusive separation. Using radio occultation experiments (not conducted by Voyager 2 at Triton), New Horizons confirmed the existence of a stratosphere (temperature increasing with height) extending to 25 km altitude at both the entry and exit locations. The entry location had a shallow troposphere (temperature decreasing with height) extending to 3.5 km altitude above the surface, while the exit location did not.