Venus before Venus Express

ARTICLE IN PRESS

Planetary and Space Science 54 (2006) 1249–1262 www.elsevier.com/locate/pss

Venus before Venus Express Fredric W. Taylor Atmospheric, Oceanic and Planetary Physics, University of Oxford, Clarendon Laboratory, Oxford OX1 3PU, England Accepted 10 April 2006 Available online 17 August 2006

Abstract An overview is given of current knowledge and mysteries about the planet Venus, with emphasis on those aspects that are intended to be studied with the Venus Express mission following orbit insertion at the planet in March 2006. r 2006 Elsevier Ltd. All rights reserved. Keywords: Venus; Atmosphere; Climate

1. Introduction Venus is the closest planet to the Earth, both in terms of distance and in terms of its physical character. The two planets are almost the same size and mean density (see Table 1), and, so far as we know, have much the same solid-body composition. The largest external differences appear in the absence of a natural satellite around Venus, the slow, retrograde rotation of the solid body of Venus, and the absence of a measurable Venusian planetary magnetic field. Internally, Venus seems to be much more volcanically active than the Earth and, perhaps as a result, has a strikingly dense and hot surface environment that far exceeds in pressure and temperature most of the pre-space age astronomers’ expectations based on models of an Earth-like planet somewhat closer to the Sun. The similarities lead us to expect that we might know the Earth better if we continue to compare it to its neighbour and twin in ever-increasing detail. The differences tell us that there are crucial aspects of solar system formation and planetary evolution, and of geology and climate physics, which we currently fail to appreciate or understand. These motivate us to carry out missions like Venus Express, which, if they are to be most effective, focus on the major unknowns and on the observed properties of Venus that are known but difficult to explain, like the high surface temperature. The goal of this paper is to summarise these Tel.: +44 1865 272903; fax: +44 1865 272924.

E-mail address: [email protected]. 0032-0633/$ - see front matter r 2006 Elsevier Ltd. All rights reserved. doi:10.1016/j.pss.2006.04.031

mysterious aspects, along with the known facts and some informed speculation about Venus, in the hope and expectation that the new mission will shed light on the mysteries, and increase and enhance our understanding of the facts.

2. The morning and evening star As the brightest of the planets, Venus has always been a much-noted feature of the night, and even the day-time, sky, always appearing within 451 of the Sun by virtue of its inferior orbit relative to the Earth. Many of the parameters listed in Tables 1 and 2 have been known from the early days of scientific observations using telescopes, including the fact that Venus has a very high reflectivity (albedo), apparently due to a thick and ubiquitous cloud cover. At 0.76, the albedo of Venus is two and a half times that of the Earth, more than offsetting the doubling of the solar constant at Venus’ mean distance from the Sun. Since Venus absorbs less radiant energy than Earth, there was no particular reason why early practitioners of what we would now call climate modelling should expect the surface temperature to be massively different from our own, and the popular vision of the surface of Venus often included oceans, deserts and steamy jungles. The Sun, with a disc twice the area it shows at the Earth, was thought to evaporate water efficiently and produce the thick and extensive cloud deck. When the composition of the atmosphere was shown by Adams and Dunham in 1934

ARTICLE IN PRESS 1250

F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

Table 1 Venus/Earth comparison (after Williams, 2005) Venus

Earth

(Venus/Earth)

Bulk parameters Mass (1024 kg) Volume (1010 km3) Equatorial radius (km) Polar radius (km) Volumetric mean radius (km) Ellipticity (polar flattening) Mean density (kg/m3) Surface gravity at equator (m/s2) Escape velocity (km/s) Bond albedo Visual geometric albedo Solar irradiance (W/m2) Equivalent blackbody temperature (K) Topographic range (km)

4.8685 92.843 6051.8 6051.8 6051.8 0.000 5243 8.87 10.36 0.76 0.65 2613.9 231.7 15

5.9736 108.321 6378.1 6356.8 6371.0 0.00335 5515 9.80 11.19 0.30 0.367 1367.6 254.3 20

0.815 0.857 0.949 0.952 0.950 0.0 0.951 0.905 0.926 2.53 1.77 1.911 0.911 0.750

Orbital parameters Semi major axis (106 km) Sidereal orbit period (days) Tropical orbit period (days) Perihelion (106 km) Aphelion (106 km) Synodic period (days) Mean orbital velocity (km/s) Max. orbital velocity(km/s) Min. orbital velocity (km/s) Orbit inclination (deg.) Orbit eccentricity Sidereal rotation period (h) Length of day (h) Obliquity (deg.)

108.21 224.701 224.695 107.48 108.94 583.92 35.02 35.26 34.79 3.39 0.0067 5832.5 2802.0 177.36

149.60 365.256 365.242 147.09 152.10 — 29.78 30.29 29.29 0.00 0.0167 23.9345 24.0000 23.45

0.723 0.615 0.615 0.731 0.716 — 1.176 1.164 1.188 — 0.401 243.686 116.750 (0.113)

Table 2 Observational parameters Distance from Earth Minimum (106 km) Maximum (106 km)

38.2 261.0

Apparent diameter from Earth Maximum (seconds of arc) Minimum (seconds of arc)

66.0 9.7

Maximum visual magnitude


4.6

Mean values at inferior conjunction with Earth Distance from Earth (106 km) Apparent diameter (seconds of arc)

41.44 60.2

to be mainly carbon dioxide, soda water oceans became the vogue for Venus. In the 1950s, it became possible to estimate the surface temperature of Venus for the first time using radio telescopes to measure the intensity of emitted microwave radiation. At wavelengths of a few centimetres, photons emitted from the surface of the planet pass almost unaffected through the cloud layers, and can be measured on Earth. The early results for Venus showed temperatures of around 400 1C, much too hot for free water or plant life.

The first space mission to Venus, Mariner 2, carried a small microwave radiometer to confirm this measurement from close range and show, by observing limb darkening, that the intense radiation was indeed coming from the surface, and not from a non-thermal source like the ionosphere. Later, the Soviet ‘Venera’ series of spacecraft made the first landings on the planet’s surface, confirming that the temperature was around 730 K, and accompanied by a pressure of nearly a hundred Earth atmospheres. At visible wavelengths, the cloud cover on Venus is complete and impenetrable, with no markings which could be associated with continents, oceans or any of the surface features which abound on the other inner planets. Instead, only extremely subtle and ephemeral markings, and some ‘scalloping’ of the terminator which separates the day and nightsides, have been reported by visual observers. Through an ultraviolet filter, like that used in the television cameras on Mariner 10, which observed Venus from a distance of 10,000 km in 1973, subtle dark markings appear in the clouds (Fig. 1). In the mid-1980s, it was discovered that much more striking contrasts can be observed at certain wavelengths in the near-infrared (IR) part of the spectrum. These are also due to the clouds, but at considerably greater depths, where large-scale meteorological activity apparently organises the clouds into patterns,

ARTICLE IN PRESS F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

Fig. 1. An image of Venus through an ultraviolet filter obtained by the Pioneer Venus Orbiter spacecraft in 1979. The contrasts in the clouds have been exaggerated by processing, showing the quasi-permanent ‘sideways Y-shaped’ feature clearly (NASA).

rather as on the Earth but in a different and still mysterious way. The discovery by D. Allen and colleagues in 1984 of near-IR emission from the nightside of Venus provided a powerful new technique for studying the lower atmosphere and surface (see Taylor et al., 1997, for a review). This emission, which is most intense within spectral ‘windows’ between strong molecular absorption bands in the 0.9–2.5 mm wavelength region, is only detected on the nightside, where it is not overwhelmed by the more intense solar flux reflected from the clouds. Near-IR imaging and spectroscopic observations of this emission by groundbased and spacecraft instruments have been used to investigate cloud particle sizes and optical thickness, the winds within the middle and lower cloud decks, and the abundances of several important trace gases, including water vapour, halides, carbon monoxide, sulphur dioxide, and carbonyl sulphide. They have provided new information about the near-surface temperature lapse rate, and the deuterium-to-hydrogen ratio. Venus Express will be the first mission with optical and spectroscopic experiments planned around the exploitation of these windows, with the prospect they offer for further progress in exploring an important and largely mysterious planetary regime which previously was accessible only to microwave sounding or entry probes. During the 1991 Venus fly-by of the Galileo spacecraft, en route for Jupiter, it was found that the principal topographical features on the surface of the planet could also be discerned in images obtained at very near-IR wavelengths. These contrasts originate in the temperature lapse rate of the atmosphere, which causes high features on Venus to appear dark in maps of the thermal emission from

1251

the surface. The very high scattering albedo of the cloud droplets in the near-IR means that the emission can diffuse to space in the spectral windows between the absorption bands of the main atmospheric constituents. Spectroscopy in these windows allows the abundance of interesting minor constituents of the atmosphere near the surface, like water vapour and carbon monoxide, to be mapped. Before the first mission to carry a surface-imaging radar to Venus flew in 1978, progress had been made in obtaining radar pictures of Venus using the large dishes at Goldstone and Arecibo to transmit and receive the pulses. Regions of high radar reflectivity appear bright, indicative of variations in either the composition or the surface morphology. In other words, bright areas can appear so either because they are composed of material which is an intrinsically better reflector, or because the surface is smoother than surrounding darker areas, or because the surface is tilted to be more nearly normal to the incoming beam than the local horizontal. The Pioneer Venus Orbiter was deployed in a high inclination orbit and so was able to map a much greater fraction of Venus’s surface than is possible when observing from the Earth, as well as obtaining better spatial resolution and relative height information. Maps were slowly built up strip by strip as the orbit precessed around the globe, taking a Venus year of 243 days to cover it completely. Veneras 15 and 16 in 1984 also mapped Venus using radar and improved the spatial resolution to around 1 km. The latest NASA mission to Venus, and the last by any agency before Venus Express, was Magellan, which operated from 1990 to 1994 and was also dedicated to radar imaging. With a resolution of 75–120 m, Magellan data produced an explosion of knowledge about the surface of Venus, taking it from the least to one of the best-explored terrains in the solar system. Of course, seeing something is only the first step towards understanding it, and the radar maps pose as many questions as they answer (Fig. 2).

3. Surface and interior The Venera landers obtained photographs of the terrain near the landing sites, revealing a sterile, scorched desert dominated by the boulders that appear strewn about the landscape. Some of the boulders have dark bands and others a patchy appearance; many have sharp edges, apparently the result of fracturing by some geological process. Most likely, the rocks are rubble from the breakup of the ejecta and lava flows associated with volcanic activity, but what process achieved this fracturing is mysterious, since running water, large daily or seasonal temperature changes and wind erosion are not available to weather rocks on Venus as they do on Earth. Some of the rocks do show evidence of erosion, possibly due to chemical action by acidic vapours in the atmosphere, or to the melting of volatile components of the rocks.

ARTICLE IN PRESS 1252

F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

The rocks rest in one view on a flat, rocky plain and in another on a mottled layer of soil or gravel (Fig. 3), giving an overall impression rather reminiscent of alluvial material, that is, deposited and modified by flowing liquid. Considering conditions on Venus, it is virtually impossible to say what processes are at work, or over what timescale,

Fig. 2. Radar map of the surface of Venus obtained by the Magellan mission in the early 1990s. In this view the N pole is near the centre, the high (red) feature just below centre is the large continent Aphrodite Terra, containing the 12-km high mountain Maxwell (white). The blue regions are depressions, possibly including primordial ocean basins. The largest of these visible here (above and to the right of centre) is Atalanta Planum. No evidence has been found for plate tectonics, leaving volcanism as the most likely process to have shaped the surface of present-day Venus.

from photographs alone. The Venera landers carried g-ray spectrometers that found uranium, thorium and potassium in the surface rocks in proportions that are consistent with a composition like terrestrial basalt. The density estimates from g-ray backscattering of 2.7–2.9 g/cm3 support this conclusion. It appears that Venus formed in a manner similar to the Moon, Earth and Mars, condensing from a molten protoplanet into shells, with the most fusible basaltic minerals making up the crust. Radar altimetry, first from Pioneer Venus then with higher resolution and coverage from Magellan, revealed a surface that is about 70% smoothly rolling plains, with about 20% lowland regions. The remaining 10% corresponds to the principal highland areas or ‘terrae’ of Venus: Ishtar, Lada, and Aphrodite, and the adjacent Beta, Phoebe and Themis regions making up a fourth major continent. In addition, there is Lakshmi Planum, a 3–4 km high plateau, bordered by mountainous ridges. Ishtar covers an area comparable to Australia and rises steeply from the surrounding plains at about 701N. The western part is a high plateau (3 km above the mean radius of Venus) bordered by tall mountains that reach a further 3 km in altitude. In the middle of Ishtar stand the Maxwell Montes, which at 11 km are high enough to tower above Earth’s highest mountain, Everest, and are much more steeply sided. The existence of such a steep and massive mountainous feature on Venus implies processes deep within the crust of the planet that created the mountain, and that continue to support its massive bulk. On Earth such a feature would most likely result from an energetic collision between surface plates, trying to move sideways into each other, in which case Maxwell would be analogous to the Himalayas on Earth. On Venus there is little other evidence for plate tectonics and it seems more likely that Maxwell is supported by vigorous convection in a large ‘hot spot’ in the crust. Gravity field experiments

Fig. 3. Pictures of the surface of Venus obtained by Venera 14 at 131N, near the eastern flank of Phoebe Regio, on 5 March 1982 (top) and Venera 9 at 321N, inside the large continent of Beta Regio, on 22 October 1975 (below). The Venera 14 site shows flat, basaltic rocks probably formed by the geologically recent break-up of volcanic lava flows, while Venera 9 shows what seem to be older, more weathered rocks sitting on a bed of finer material, like sand or soil. The details and ages of the processes that shaped these localities remain unknown, as does the stratigraphy and composition of the layers underneath the exposed surface. The chemical interaction of surface materials with the atmosphere may have a key role in explaining the extreme climate on Venus today.

ARTICLE IN PRESS F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

and bi-static sounding of the more mountainous regions by Venus Express may shed further light. The reason Venus apparently does not have continentsized plates, moving relative to each other as they do on the Earth, is not known. Possibly, the low water content or some other aspect of the composition of the crust, its thickness, and the high surface temperature, cause it to crack more easily than on Earth, so major plates break up instead of moving around. Narrow linear structures (chasmas) and tesserae or tiled terrains can be seen all over the surface of Venus, and these and other tectonic features (including the montes) may be products of the cracking of the crust that occurs in response to convection in the mantle below. Aphrodite Terra is the second most prominent highland region on Venus, with an area about equal to that of Africa in an elongated shape that has been likened to a scorpion stretching for about 10,000 km along and south of the equator. The highlands at the eastern end contain steep valleys on an enormous scale not found on Earth. The largest, Diana Chasma, is comparable to the Valles Marineris on Mars. The lowlands (Planitiae) are fairly featureless, probably a result of having been flooded by lava relatively recently, in some cases several times with fresh flows partially covering the earlier ones. Because of this resurfacing, much of the surface of Venus appears geologically very young. Dating the surface by the usual method involving the study of impact craters is made difficult by the thick atmosphere, which prevents the smaller and more frequent impactors from making an observable impression of the surface. The craters that have been seen, several hundred in number, are quite pristine in appearance and uniformly distributed over Venus, suggesting that the renewal of the surface has been geologically recent and planet-wide. The Magellan images show that most of the surface of Venus is made up of features attributable to volcanic activity. The largest volcanoes are the shields, found mostly in high regions lying 3–5 km above the surrounding area and featuring lava flows which often extend for hundreds of kilometres from the central caldera. The smaller volcanoes are called anemones, ticks, and arachnids, depending on their appearance. Anemones are characterised by ‘hairy’ flow patterns typically 50 km across, radiating out from a central source of magma. Ticks are flat, circular volcanic domes about 25 km in diameter flanked by strongly defined radial ridges and troughs. The arachnids are volcanic mounds which appear to have collapsed, cracking the crust and producing a distinctive insect-like shape. Similar to ticks but without the ‘legs’ are the steep-sided, flat-topped volcanoes known as pancake domes. The shape suggests that the lava that formed them was of higher viscosity than that emanating from the more Earth-like large volcanoes, and so did not flow as freely or as far. The nearest terrestrial analogues to pancake domes are on the sea bed, where the surrounding high-density fluid affects the cooling and solidification of the dome.

1253

There are remarkable river-like features on Venus, some of them extending thousands of kilometres from their volcanic sources to the lava-filled flood plains. Several have cut deep, meandering valleys, a process that must have required large volumes of liquid to flow over long periods of time. The fluid involved is obviously not water, but something with a melting point that is not too different from the mean surface temperature on Venus. Assuming this temperature has not varied substantially since the valleys formed, the most likely candidate would be a carbonate-rich mineral such as carbonatite, although there are others that cannot be ruled out, even including the low melting-point metals like lead or tin. Although the plains clearly are solid, over most of their area at least, we do not know for certain that these rivers, if they are different in composition from the large-scale flows, are not still running in places. Another curiosity is that the higher features on Venus seem to be not only relatively unweathered, but actually plated with something, which gives everything more than about 2 km above the mean surface height an unnaturally high radar brightness. Early guesses as to the composition of these ‘snowcaps’ were made by looking for something which condenses at the temperature and pressure of the observed lower altitude boundary. They included the metal tellurium, or the compound iron pyrites (FeS), also known as fool’s gold from its resemblance to the precious metal. More recent work by Schaefer and Fegley (2004), based on chemical equilibrium arguments, favours a mixture of lead and bismuth sulphides. 4. Atmospheric composition and clouds A comprehensive and still relevant summary of present knowledge about the composition of Venus’s atmosphere, plus an account of how the data were obtained, has been given by von Zahn et al. (1983)); an update is presented by de Berg et al. (this issue). The abundances derive mainly from mass spectrometer and gas chromatograph measurements made on the Pioneer Venus and Venera descent probes. A summary is included in Table 3. Several of the minor constituents exhibit striking amounts of temporal and spatial variability, indirectly revealing major characteristics of the planet and its atmospheric circulation and meteorology. Sulphur dioxide was observed in ultraviolet measurements made by Pioneer Venus to show large variations in its abundance near the cloud tops, which Esposito (1984) interpreted as evidence for active volcanism at the surface. During the Galileo flyby in 1991, near-IR measurements revealed an equator-topole gradient in the abundance of tropospheric carbon monoxide (Collard et al., 1993), which Taylor (1995) suggested could be characteristic of a hemispherical Hadley circulation that extended from the lower thermosphere at around 120 km all the way down to the surface. Finally, water vapour measurements above, below and within the cloud layers show a baffling disparity that is presumably

ARTICLE IN PRESS 1254

F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

Table 3 Atmosphere Surface pressure Surface density Scale height Total mass of atmosphere Average temperature Diurnal temperature range Wind speeds Mean molecular weight Atmospheric composition (near surface, by volume) Major Minor (ppm)

92 bar 65 kg/m3 15.9 km 4.8  1020 kg 737 K (464 1C) 0 0.3–1.0 m/s (surface) 43.45 g/mol

Carbon dioxide (CO2) 96.5% Nitrogen (N2) 3.5% Sulphur dioxide (SO2) 150 Argon (Ar) 70 Water (H2O) 20 Carbon monoxide (CO) 17 Helium (He) 12; Neon (Ne) 7

linked both to cloud formation and dissipation processes and meteorological activity in Venus’ atmosphere (Taylor, 2006a, b). Observations of Venus at visible wavelengths, even from close-up, show little or no detail in the clouds, which appear to the optical astronomer observing from the Earth to form a uniform blanket over the whole planet. However, it was first noted in the 1920s that photographs taken through an ultraviolet filter show blotchy dark features in the cloud. Under the best observing conditions, the dominant pattern in these markings shows a characteristic shape, like a letter ‘Y’ laid sideways (Fig. 1). The contrasts are due to some ultraviolet-absorbing substance that is non-uniformly dispersed through the clouds. Sulphur dioxide behaves in this way, and is definitely present in spectroscopic observations, but its spectrum does not match that of Venus precisely at all wavelengths. Some other material, probably another sulphur compound or even one of the allotropes of elemental sulphur, which also absorbs ultraviolet more than visible radiation, must be contributing also. The main cloud deck on Venus extends from about 45 to about 65 km above the surface, with haze layers above and below (Fig. 4). Within this gross structure, detailed layering occurs and particles of different sizes congregate at different height levels. The particles range in diameter from less than 1 to over 30 mm and tend to a trimodal size distribution, with the commonest diameters falling towards the ends of the overall range and in the 2–3 mm region. Spectroscopic, polarimetric and other evidence yields a composition of 75% H2SO4 and 25% H2O for the intermediate size or ‘mode 2’ droplets. The composition of the smaller, ‘mode 1’ drops is unknown; these form an aerosol haze extending throughout the cloud layer. Most of the mass of the clouds is in the big ‘mode 3’ drops; these may be more evolved sulphuric acid drops constituting a tail to the mode 2 distribution

Fig. 4. Temperature (black curves, bottom scale) and cloud density (grey curve, top scale) profiles for Venus, based on measurements from several different instruments on the Pioneer Venus orbiter and entry probes. The global thermal structure within and beneath the clouds has been sampled but not mapped in enough detail to understand Venusian meteorology. Only tentative theories for the extreme greenhouse warming exist, and the detailed structure, variability and compositional variations in the clouds are also unknown.

curve, or (particularly since there is inconclusive evidence of a non-spherical shape) a separate mode altogether with a different, unknown, composition. The sulphuric acid droplets are probably formed when H2O and SO2 (the latter presumably of volcanic origin) combine photochemically near the cloud top level. It is difficult to explain the details of the size distribution, particularly to explain the existence of particular modes, and their multiplicity. Compositional contrasts and dynamical effects may be at work, but at present the observations which would elucidate these are in short supply. In spite of their physical depth, the clouds are not completely opaque at all wavelengths, because they are quite efficient scatterers in the visible and near-IR regions of the spectrum. The level of illumination in the Venera pictures was higher than had been expected. Even with the Sun 601 above the horizon, it was thought that the thick clouds would prevent more than a trace of sunlight from reaching the ground; instead a light level which has been compared to that on Earth during a thunderstorm was discovered and the searchlights which the spacecraft carried were unnecessary. In 1978, radiometers on the Pioneer Venus probes measured the solar flux and found that 2.5% of the total falling on the planet actually reaches the surface.1 A corollary of this is that the hot surface and lower atmosphere of Venus, which emit strongly in the IR, can be seen from outside the atmosphere through the clouds in observations made at wavelengths outside regions 1 The estimates from the Soviet landers, up to and including Veneras 11 and 12, was somewhat higher, at 3.457.45% (Moroz et al., 1983).

ARTICLE IN PRESS F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

of strong CO2 or H2O absorption. The best study of this phenomenon, discovered only in the 1980s when the first phase of dedicated Venus exploration by spacecraft was essentially over, were made by the Near-Infrared Mapping Spectrometer on the Galileo Jupiter Orbiter during its intermediate flyby of Venus in 1990. Maps of Venus obtained by this instrument show large horizontal variations in the thickness of the main cloud deck in considerable detail (see Fig. 5). The large-scale structure of the clouds, in the horizontal and in the vertical, must be due to a variable balance between dynamical transport and the production and loss processes for the cloud materials. Except that the sharp lower boundary of the entire cloud deck is probably due to evaporation beyond a high temperature threshold, and that sulphuric acid production is to be expected at the cloud tops, via the reactions: CO2 þ SO2 þ hn ! CO þ SO3 ; SO3 þ H2 O ! H2 SO4 : These processes remain largely mysterious and require further investigation. In addition to the general structure of the various layers, specific puzzles are the nature of the absorber that produces the contrasts seen in UV-blue

1255

images of the planet, and the composition of the large ‘mode 3’ particles seen in the Pioneer Venus Large Probe data. Indirect evidence suggests that these are likely to be solid, and therefore not H2O, H2SO4 or even HCl. Another feature of the Venusian clouds which has been hotly debated is the question of whether or not lightning is present. On theoretical grounds, this was thought rather unlikely, because the clouds are too tenuous, although localised storms and clouds of volcanic ejecta could provide the right conditions. The Galileo observations from February 1990 (Gurnett et al., 1991) provided the hardest experimental evidence to date by detecting impulsive radio signals in the 100 KHz–5.6 MHz frequency range, for which lightning is the only known source. Lightning was not detected optically, however, in spite of a search by the Galileo imaging experiment (Belton et al., 1991). Finally, the question of CO2 clouds on Venus is worth a mention, as an interesting curiosity more than as a crucial aspect of the climate and meteorology of the planet. The mean temperature of the atmosphere above the thin H2SO4 cloud, at altitudes of around 90 km, is in the region of 180 K (Fig. 3). During the night, the predominantly CO2 atmosphere radiates strongly to space and temperatures probably drop below 160 K, cold enough for CO2 ice clouds to form. The Pioneer Venus Orbiter Infrared Radiometer noted anomalously strong limb darkening near the dawn terminator, between 601 and 951 solar longitude and up to 151 latitude, that Taylor (1981) tentatively attributed to carbon dioxide (or possibly water ice) cloud formation near the 90 km level. The much more sophisticated and extensive observations expected from Venus Express and the Japanese Planet-C missions should readily settle this question. 5. Thermal structure and energy balance

Fig. 5. An image of the nightside of Venus in the near-infrared ‘window’ at a wavelength of 2.3 mm, obtained by the Near-Infrared Mapping Spectrometer aboard the Galileo spacecraft during its flyby in February 1990 (Carlson et al., 1991). The details, expressed in false colour, are due to the lower-level cloud structure, back-lit by thermal emission from the hot lower atmosphere and surface. The clouds are seen to be highly nonuniform, leading to brightness variations which range over more than an order of magnitude from white and red (thin cloud regions) to black and blue, representing thick clouds. This indicates surprisingly active meteorological behaviour in the deep atmosphere, where long radiative and dynamical time constants might be expected to suppress such activity.

Enough sunlight diffuses through the cloud layers to provide about 17 W/cm2 of surface insolation (Tomasko et al., 1983). This is about 10% of the total absorbed by Venus as a whole, including the atmosphere, assuming a fraction A ¼ 0:76 is reflected from the planet without being absorbed. Thus heated at and near the surface, the lower atmosphere forms a deep convective region, the troposphere. Within the troposphere the atmosphere cannot cool significantly by radiation to space, because the opacity of the overlying layers is large. The greenhouse effect, whereby short-wavelength solar radiation heats the lower atmosphere more easily than the longer thermal wavelengths can cool it, raises the surface temperature significantly above that which would apply on an airless planet. The effect is particularly extreme for Venus, where the surface temperature must rise to 730 K in order to force enough IR cooling to balance the incoming solar energy. An airless body with the same albedo and heliocentric distance as Venus would reach equilibrium for a mean surface temperature of only about 230 K. Nevertheless,

ARTICLE IN PRESS 1256

F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

Fig. 6. Illustrating some of the processes though to be involved in producing the climate on Venus, which is much more extreme, compared to the Earth, than anyone expected prior to the first space missions to the planet. It remains a major task to explain how these, and possibly other, factors combine to give Venus its high surface temperature and pressure, and how that may evolve in the past. Another key question is how the climate will change in the future, when the volcanoes that stud the surface of Venus, some of them probably very active at present, finally subside. A much more Earth-like regime is one possibility suggested by models (Taylor, 2006a, b), but there remain many uncertainties.

measurements by the Pioneer Venus Orbiter of the net IR emission and the total reflected solar energy (Schofield and Taylor, 1982) confirmed that the planet is in overall balance to within the accuracy of the measurement. The Pioneer and Venera measurements also showed that the net thermal emission from the planet is a much weaker function of latitude than it is on the Earth, despite a distribution of incoming energy from the Sun that is more biased towards the equatorial region by the small obliquity of Venus. The basic reason is undoubtedly that the atmosphere on Venus is an efficient transporter of sensible and latent heat from equator to poles, not only because it is dense and deep and therefore has a high thermal capacity, but also because it apparently has a less turbulent regime that allows advection to proceed more freely. Such general assumptions rest on only very fragmentary evidence, however, and the processes actually at work have yet to be fully revealed, modelled or understood. The new capabilities afforded by near-IR sounding of the deep atmosphere, to be fully exploited for the first time by the spectroscopic and imaging instruments on Venus Express, are likely to make great strides in this area. A detailed understanding of the greenhouse effect on Venus, and of why it finds equilibrium at such an unexpectedly high surface temperature, requires new insights into the composition of the atmosphere and the

cloud layers, the radiative behaviour of the gases under extremes of temperature and pressure, cloud chemistry and microphysics, the role of volcanism, and the chemical equilibrium between the atmosphere and surface minerals (Fig. 6). Reactions between atmospheric carbon dioxide and carbonate and silicate minerals in the crust may hold the key to understanding the high surface temperature and pressure, as discussed for instance by Bullock and Grinspoon (1996), following the original suggestion of Urey (1952). Despite the fact that reactions such as CaCO3 þ SiO2 2CaSiO3 þ CO2 reach equilibrium with CO2 at temperatures and pressures typical of those found on Venus, and would therefore seem to offer a solution to the dilemma of Venus’ extreme climate, there are many practical questions that remain unresolved about how such a mechanism would operate in practice (see, for instance, Hashimoto and Abe, 2005). Convection in the troposphere carries energy upwards to the base of the stratosphere, where radiative cooling to space can occur strongly. On Venus, this level (the tropopause) occurs about 60 km above the surface (Fig. 3). Above the troposphere, in the middle atmosphere or mesosphere, the temperature tends to a constant value with height, because the atmosphere here is optically thin. To a first approximation, each layer tends to find the same

ARTICLE IN PRESS F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

equilibrium temperature, determined by the balance between the absorption of upwelling IR from the surface and troposphere and cooling to space, if no significant absorption of direct solar energy takes place. Overall, the mean vertical temperature profile on Venus follows these general expectations and is therefore fairly unremarkable. The superimposed details observed by thermometers on entry probes and by IR remote sensing from orbiters reveals a great deal of interesting detail, however. These include (i) discontinuities in the vertical lapse rates and hence the static stability of the atmosphere, in distinct slabs that correlate between different probes; (ii) a vertically propagating, solar fixed thermal tide that is predominantly wavenumber 2 in character, accompanied by a variation of about 1.5 km during the solar day in mean cloud-top height, dominated in this case by the wavenumber-1 component; (iii) strong latitudinal variations in the temperature structure in the middle atmosphere, 60–100 km, that show a steady increase from equator to pole. Dynamical theories have been advanced to explain each of these, in the form of (i) multiple, stacked Hadleytype circulation cells; (ii) a variant on terrestrial tidal theory that shows wavenumber 2 propagates more strongly than wavenumber 1 under model Venus conditions and (iii) pressure gradients associated with the vertical superrotation and corresponding thermal gradients that are far from radiative equilibrium. All of these are exciting ideas that require considerable elaboration and testing with new and more detailed data, whereupon much can be learned about meteorological processes that, like many aspects of the Venusian environment, have elements of both the Earth-like and the bizarre. Above the mesopause the thermosphere begins, a lowdensity region that takes its name from the fact that temperature increases with height on the dayside due to the absorption of mainly ultraviolet photons from the Sun, principally in the extreme ultraviolet portion of the spectrum. Energetic particles in the solar wind also contribute, and the actual thermospheric temperatures vary considerably with solar activity and sunspot cycle. The thermosphere of Venus is cooler than Earth’s, because of the greater abundance of carbon dioxide, which is very efficient at radiating heat to space. Above about 150 km, the temperature is approximately constant with height on the dayside at about 300 K. The terrestrial thermosphere is the seat of rapid winds, up to 1000 m/s or more, and this tends to redistribute energy originally absorbed from the Sun over the dark as well as the sunlit hemisphere. The result is a modest day–night difference of around 200 K about a mean temperature of 1000 K. On Venus however, the night-time temperature in the thermosphere is very low, around 100 K. The transition from the day to nightside values of temperature on Venus also show surprisingly steep gradients (Keating et al., 1979) and modellers have great difficulty in reproducing both the minimum temperature and the short distance across the terminator with which it is attained. Somehow the dynamics of Venus’s

1257

thermosphere is such that the flow of air in response to the temperature gradient is inhibited. 6. Atmospheric dynamics It was known by the early 1960s that the Y-feature seen in ultraviolet images rotates around the planet in a period of only 4–5 days. This implied wind velocities of 100 m/s at the cloud tops, surprisingly rapid for such a slowly rotating planet. The solid surface of Venus rotates at only about 2 m/s, or once every 243 days. Comparison of the Y-shaped markings in a composite of Earth-based observations in 1966, a mosaic of Mariner 10 ultraviolet photographs taken during the period from 5–12 February 1974, and as seen in similar pictures obtained by Galileo, when it flew past Venus en route to Jupiter in February 1990, all show a striking similarity. Evidently the dynamical processes in the atmosphere of Venus give rise to stable wave modes that are reflected in the cloud patterns, but when it comes to the details, the meteorology of Venus is as mysterious as its surface geology. Careful measurements of the propagation velocities of small scale features that move with the winds and that of the large Y (representing the phase velocity of the wave or waves, superimposed on the wind velocity) confirm that the bulk velocity of the atmosphere in the zonal direction (parallel to the equator) at low latitudes is over 100 m/s, a result which is borne out by Earth-based measurements of the Doppler shifting of spectral lines. The global-scale waves propagate upstream at about 20–30 m/s. Both of these velocities are much larger than the apparent velocity of the Sun with respect to an observer on the surface. Probe and remote sounding measurements show that the rate at which the atmosphere circulates around the equatorial regions varies considerably with height, its maximum speed occurring near the cloud tops where the ultraviolet markings apparently originate. Above the clouds, the 100 m/s winds are decelerated sharply by the pressure gradient which results from the temperature distribution at those levels. The remarkable discovery that the air temperature is some 15–20 1C warmer at the pole than the equator from 70–95 km above the surface relates to this: dynamical models imply that the pressure gradient implied by this temperature structure is sufficient to arrest the zonal winds completely by 85–90 km altitude if cyclostrophic balance is assumed. Below the clouds, the winds fall gradually in velocity as the atmosphere becomes denser. Doppler tracking of the Pioneer probes shows a wind speed of less than 10 m/s at the 10 km altitude level and close to zero at the surface (Fig. 7). All of the zonal winds are westward (in the same direction as the rotation of the planet), suggesting that angular momentum is being delivered to the atmosphere by the solid body of the planet and transported upwards. Alternatively, it has been proposed that the Sun exerts a torque on the atmosphere and so supplies external angular momentum. This it certainly will do, since the density of

ARTICLE IN PRESS 1258

F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

Fig. 7. Venus wind profiles from the Pioneer Probes. The zonal winds are retrograde and reach speeds in excess of 100 m/s around the equator at cloud-top heights. Tracking of cloud features in both ultraviolet and near infrared images has confirmed that there is a small net equator-to-pole velocity component of a few m/s both at the visible cloud tops at around 60 km altitude and in the main cloud layer some 15 km lower. The processes that accelerate the zonal wind to such high values (more than 50 times the rotation rate of the solid planet, at relatively high atmospheric densities) are still qualitatively unknown.

the atmosphere is non-uniformly distributed with solar longitude (local time of day) because of thermal tides induced by solar heating. In fact, the semidiurnal component of the tide, on which the torque principally is exerted, has been observed to be unexpectedly large, relative to the diurnal component, on Venus, which favours this mechanism. Whether the effect is large enough to accelerate the atmosphere to the speeds observed is a subject of ongoing debate. There is even a possibility that the slow retrograde rotation of the planet itself may have been established, over geological time, by the torque which the atmosphere exerts on the planet—the reverse of the earlier theory. Less dramatic than the zonal winds, but of even greater significance for the general circulation of Venus, is the observed migration of the ultraviolet markings away from the equator towards both poles (i.e., in the meridional direction) at speeds of less than 10 m/s. Galileo/NIMS also observed poleward motions in the deeper cloud structure, imaged in the near-IR windows. The general impression is of two gigantic circulation cells, one to each hemisphere, in which (heated) air rises at the equator and (cooled) air descends at the poles, travelling more or less horizontally in

between, polewards above the clouds and equatorwards below. Such a flow is the characteristic of a Hadley cell, the simplest circulation regime which can occur in a planetary atmosphere. This kind of structure was proposed for the Earth’s atmosphere by Hadley as long ago as 1735. It seemed logical to him that rising air at the warm equatorial regions and falling air at the cool poles, would lead to poleto-equator flow near the surface and to motion in the opposite direction at higher levels. This does in fact happen, but in our own atmosphere this simple structure is greatly modified by the development of baroclinic instabilities in the motion, and the smooth flow of the Hadley regime tends to break down under the influence of the Earth’s spin. The net result on Earth is that Hadley’s cell extends from the equator only to mid-latitudes, with other, smaller, cells taking over the transport nearer the Pole. On slowly rotating Venus, it appears that the basic Hadley configuration exists in a less modified state, spanning each hemisphere. Dramatic evidence for this type of circulation was obtained by IR measurements from the Pioneer Venus orbiter, which provided the first observations from above the polar regions. Greatly enhanced amounts of IR flux were found to be emerging poleward of 801N latitude, in a localised, elliptical region which evidently is a clearing in the cloud cover forced by the descending air in the return branch of the Hadley cell (Fig. 8). This clarification of the circulation regime existing near the cloud tops also gives us some clues to other fundamental questions about Venus, namely its global cloud cover and its high surface temperature. Probably, on a long-term average, the air is rising slowly everywhere on Venus and descending only near the poles, with the cloud-top patterns seen in the ultraviolet pictures and the deep atmosphere cumulus dynamics seen in the near-IR pictures superimposed as the short-term disturbances or weather. The resulting planetwide blanket of cloud plays a key role in trapping heat radiation in the lower atmosphere and raising the surface temperature. As already noted, sulphuric acid droplets are better than water at doing this, because their optical properties are such that they scatter radiation conservatively at ultraviolet, visible, and near-IR wavelengths, where the Sun emits most of its energy, and absorb in the middle and far IR, where Venus cools by thermal emission to space. Near-IR spectroscopic measurements in atmospheric ‘windows’, that is, wavelength regions where the main atmospheric gases are weakly absorbing, penetrate the clouds, in some windows all the way to the surface. As shown in Fig. 5, this type of observation reveals the global cloud morphology at depth, not just in the cloud-top region, although they are possible only on the nightside of the planet, because the emitted signals at these short IR wavelengths are dominated by reflected sunlight on the dayside. They show that the optical thickness of the clouds, far from being largely uniform as had been supposed, is very variable, by factors of 10 or more. Regions of thick

ARTICLE IN PRESS F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

1259

Fig. 8. The N polar vortex on Venus, in UV light from Mariner 10 (top left) and at a wavelength of 11.5 mm in the thermal infrared from Pioneer Venus Orbiter (Taylor et al., 1980). The latter show a ‘snapshot’ of the vortex (top right), a 72-day average, showing the cold polar collar, and a similar average in a coordinate frame rotating every 2.7 days, showing the dipole structure. The existence of a vortex is not surprising, but its detailed collar-dipole structure has so far defied explanation and modelling.

and (relatively) thin clouds form patterns suggestive of large-scale cumulus dynamics, presumably with the cloud material actively condensing and dissipating in rising and falling air associated with weather systems, although the details are lacking because of a shortage of high-resolution data in space and time. Galileo, which made only a very brief encounter with Venus on its way to Jupiter in 1990, has been the only spacecraft equipped for near-IR imagery to visit Venus since ground-based observers in the mid1980s discovered the existence of the spectral windows. Venus Express, especially the VIRTIS and VMC instruments, will provide greatly extended coverage in space, time and wavelength, not only in the equatorial regions but also over the poles, and in the collar and dipole regimes. Some of the other important wave phenomena seen on Venus are the circumequatorial belts and the bow-like waves. The circumequatorial belts are very narrow (less than about 50 km in width), very long (lengths of the order of thousands of kilometres) and transient. As many as five have been seen at once, evenly spaced by about 500 km and always aligned parallel to the equator. They appear in 1 or 2 h and propagate, always in a southerly direction, for

about 12
112 days at about 20 m/s (Belton et al., 1976). The most satisfactory explanation for the belts is that they are some form of gravity wave, resonances in the atmosphere caused by density variations propagating as waves under the influence of gravity as the restoring force. They are common in the Earth’s atmosphere, where temperature fluctuations, associated with the density waves, lead to condensation in the thermal troughs. Something similar may be happening on Venus, although it is far from clear in this case what is exciting the waves. It could be turbulence in the strongly heated subsolar region, or perhaps some lower atmospheric wave propagating upwards. Still, it is difficult to explain why the waves always seem to travel from the north to the south. The bow waves were named for their shape (like that of a bow, as in archery). However, it turns out that they probably have something in common with the bow waves associated with the passage of a ship across water, as well. On Venus we see powerful ‘boiling’ of the atmosphere in the region directly below the Sun, i.e., at local noon. The rising of the heated air in this subsolar disturbance, visible as convection cells in the images, interferes with the

ARTICLE IN PRESS 1260

F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

smooth, high velocity flow of the upper atmosphere and generates ‘ripples’ in the clouds. However, this explanation is even more limited than its obvious oversimplification would imply, since the waves travel downstream behind the subsolar zone, whereas oceanographic bow waves remain fixed with respect to the disturbance producing them. Two other puzzling Venusian atmospheric phenomena are observed only in IR images of the planet. These use the thermal emission from the planet as a source, and so cover the dark or unfavourably illuminated portions of the planet as well as the dayside. They also reveal the temperature and vertical structure of the phenomena that they observe. The circumpolar collar is a current of very cold air that surrounds the pole at a radial distance of about 2500 km. It is nearly 1000 km across, but only about 10 km thick in the vertical direction. Temperatures inside the collar are about 30 1C colder than at the same altitude outside, so the feature generates pressure differences that would cause it to dissipate rapidly were it not continually forced by some unknown mechanism. Inside the collar lies the region of reduced cloud cover caused by the descending branch of the Hadley cell. Because of the zonal momentum transported from lower latitudes, the descending air is also rotating rapidly, forming a polar vortex analogous to the eye of a terrestrial hurricane or whirlpool but much larger and more permanent. Interestingly, however, the eye of the Venus polar vortex is not circular but elongated, and with brightness maxima (presumably corresponding to maximum downward flow) at either end. This gives the polar atmosphere a ‘dumbbell’ appearance and has led to the name polar dipole for the feature. The dipole rotates about the pole every 2.7 Earth days, i.e. with about twice the angular velocity of the equatorial cloud markings. If angular momentum were being conserved by a parcel of air as it migrated from equator to pole the dipole might be expected to rotate five or six times faster. In fact, the ultraviolet markings are observed to keep a roughly constant angular velocity (solid body rotation) from the equator to at least 601 latitude, presumably accelerating poleward of this. This assumes, of course, that the rotation of the dipole represents the actual speed of mass motions around the pole and not simply the phase speed of a wavelike disturbance superimposed on the polar vortex. Interestingly, the thermal tide on Venus around the equatorial regions also has two maxima and two minima. (The thermal tide is simply the diurnal increase and decrease of temperature caused by the rising and setting of the Sun.) This does not seem to be connected with the polar dipole, since the two regions are separated by a narrow latitude band apparently free of planetary-scale waves, as well as by the predominantly wavenumber-one collar. The Earth’s atmosphere has a small wavenumbertwo component superposed on the familiar early-afternoon maximum to post-midnight minimum cycle, but this component dominates on Venus. For once, the dynamical theory of atmospheric tides, as developed for Earth, shows

when applied to Venus that the observed state of affairs can be explained as primarily a consequence of the long solar day on Venus (Fels et al., 1984). It is puzzling, even alarming, that the same cannot be said of most of the other major characteristics of the atmosphere of the closest and most Earth-like planet (see Fig. 9). 7. Magnetic field, space environment and atmospheric escape It is generally thought that Venus accreted and differentiated in a manner similar to its near-twin, the Earth, leading in both cases to a molten iron-rich core of about half the total radius, overlaid by a crust of mainly silicate material. The mean density of 5.25 g/cm for Venus is consistent with such a picture. It used to seem reasonable, therefore, to expect that Venus should also have an Earth-like intrinsic magnetic field, until the era of spacecraft exploration when it gradually became clear that Venus has little or none (Luhmann and Russell, 1997). The Pioneer Venus Orbiter made measurements beginning in 1979 that included repeated low-altitude passes that confirmed that the planet has negligible internal field, with an upper limit of 10
5 times that of the Earth. Such fields as were observed by Pioneer in near—Venus space were attributed to solar wind interactions directly with the planet and its atmosphere, which, at high levels, contains layers of charged particles known as the ionosphere. The upper boundary of this, the ionopause, is the surface at which the dynamic pressure of the solar wind is in balance with the thermal pressure of the ions and electrons that make up the ionosphere The height of the ionopause varies with solar activity from around 250 km to more than 800 km from the solid surface of the planet. It deflects the solar wind flow around the planet, with the formation of a bow shock several thousand km upstream (Fig. 10). While the detailed physics of the processes outlined above remain of great interest, the key question for the study of the induced magnetosphere around Venus is the role of the deflected solar wind in carrying off atoms and ions derived from atmospheric molecules, particularly the lighter elements, and especially hydrogen. Venus may once have had a massive ocean that was slowly removed by dissociation of water vapour in the upper atmosphere and subsequent loss of the hydrogen. This scenario is supported by the expectation that Venus was initially water-rich, like the Earth, and also by the strong evidence of fractionation in the isotopes of hydrogen found on Venus, where the deuterium-to-hydrogen ratio is more than 100 times that found on Earth or in meteorites. The implications for the present high surface pressure and consequent extreme climate on Venus are obvious; the mass of the current atmosphere represents a balance between emissions from the crust by volcanism, the chemical recombination of atmospheric molecules with the surface, and escape to space. The details and relative proportions of these

ARTICLE IN PRESS F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

1261

Fig. 9. A summary of the main features of the atmospheric circulation on Venus, and some of the principal questions. The latter include the nature of the polar dipole feature, the extent of the Hadley cell circulation and how it may be balanced by a counter-rotating polar cell and eddy momentum transfer, and the drive for the zonal super-rotation and its balance via the reverse temperature gradient in the stratosphere.

Fig. 10. The solar wind interacts directly with ions and electrons in the Venusian ionosphere, so no Earth-like magnetotail or trapped radiation belts are produced. However, Venus does exhibit a bow shock, a magnetosheath or ionosheath inside the bow shock, and an ionopause 250 km or more above the surface where the solar wind pressure is balanced by the ionosphere thermal energy (NASA).

processes and budgets, current and historical, remain quite unknown. The Neutral Mass Spectrometer on Pioneer Venus Orbiter observed CO2, CO, O, N2, N, NO, He and H in the upper atmosphere of Venus. The UV spectrometer on the same mission established the presence of a corona of hot atoms, mainly H, O and C, around the planet. The homopause on Venus is at an altitude of about 135 km, leading to an increasing preponderance of the lighter species above this level. However, the lightest, atomic

hydrogen and helium, are minor constituents in the atmosphere as a whole while atomic oxygen is produced in large quantities in the upper atmosphere by the photodissociation of carbon dioxide by solar UV, viz. CO2+hn-CO+O. The net effect is that O atoms are the dominant species at altitudes above 170 km, especially during the daytime, followed by atomic hydrogen, helium, and molecular hydrogen (von Zahn et al., 1983). Model calculations by Lammer et al. (2006) compare the various loss processes for these species. Thermal loss, an important process for hydrogen loss on Mars, is negligible on Venus because of the much greater mass of the latter. For the same reason, sputtering (the removal of atmospheric atoms or molecules by momentum transfer in collisions with solar wind particles) is important only to the extent that sputtered particles contribute to the hot particle coronae, where other processes occur. The dominant mechanisms for hydrogen loss are thought to be H+ ion outflow accelerated along magnetic field lines on the nightside of the planet, and for O+ the same process augmented by the formation of ionospheric plasma clouds triggered by the Kelvin–Helmholtz instability. A key question is whether the net loss rates, by all processes, for hydrogen and oxygen are in the ratio 2:1 as would be expected if the source molecule is water vapour and there are no large sinks of atmospheric oxygen on Venus’s surface. Within large uncertainties, Lammer et al. (2006) find that this may indeed be the case for Venus,

ARTICLE IN PRESS 1262

F.W. Taylor / Planetary and Space Science 54 (2006) 1249–1262

although not for Mars where the ratio is about four times larger and a surface sink for oxygen almost certainly must be invoked. The ASPERA and Magnetometer measurements planned for Venus Express are designed to further elucidate this ratio and the loss rates for hydrogen, deuterium, oxygen and other species. 8. Conclusion Despite being the closest planet to the Earth and much contemplated and explored, many aspects of Venus as a planet remain enigmatic. The most urgent problem, and the primary target for the new generation of Venus missions, must be to understand the remarkable state of the climate, and the processes that are responsible, including volcanism, the atmospheric circulation, cloud chemistry, and atmospheric escape processes. All of these, and a great many other persistent mysteries that unfolded during the first great era of Venus exploration in the two decades from 1965 to 1985, should be much clearer if the Venus Express mission succeeds in making its planned observations during the next few years. References Belton, M.J.S., Smith, R.S., Schubert, G., Del Genio, A.D., 1976. Cloud patterns, waves and convection in the Venus atmosphere. J. Atmos. Sci. 33, 1394–1417. Belton, M.J.S., Gierasch, P.J., Smith, M.D., Helfenstein, P., Schinder, P.J., Pollack, J.B., Rages, K.A., Morrison, D., Klaasen, K.P., Pilcher, C.B., 1991. Images from Galileo of the Venus cloud deck. Science 253, 1531–1536. Bullock, M.A., Grinspoon, D.G., 1996. The Stability of Climate on Venus. J. Geophys. Res. 101, 7521–7529. Carlson, R.W., Taylor, F.W., et al., 1991. Galileo infrared imaging spectroscopy measurements at Venus. Science 253, 1541–1548. Collard, A.D., Taylor, F.W., Calcutt, S.B., Carlson, R.W., Kamp, L., Baines, K., Encrenaz, Th., Drossart, P., Lellouch, E., Be´zard, B., 1993. Latitudinal distribution of carbon monoxide in the deep atmosphere of Venus. Planet. Space Sci. 41 (7), 487–494. Esposito, L., 1984. Sulphur dioxide: episodic injection shows evidence for active Venus volcanism. Science 223, 1072–1074. Fels, S.B., Schofield, J.T., Crisp, D., 1984. Observations and theory of the solar semidiurnal tide in the mesosphere of Venus. Nature 312 (5993), 431–434.

Gurnett, D.A., Kurth, W.S., Roux, A., Gendrin, R., Kennel, C.F., Bolton, S.J., 1991. Lightning and plasma wave observations from the Galileo flyby of Venus. Science 253, 1522. Hashimoto, G.L., Abe, Y., 2005. Climate control on Venus: comparison of the carbonate and pyrite models. Planet. Space Sci. 53, 839–848. Keating, G.M., Taylor, F.W., Nicholson, J.Y., Hinson, E.W., 1979. Short term cyclic variations of the Venus upper atmosphere. Science 205, 62–65. Lammer, H., Lichtenegger, H.I.M., Biernat, H.K., Erkaev, N.V., Arshukova, I., Kolb, C., Gunell, H., Lukyanov, A., Holmstrom, M., Barabash, S., Zhang, T.L., Baumjohann, W., 2006. Loss of hydrogen and oxygen from the upper atmosphere of Venus. Space Sci. Rev., in press. Luhmann, J.G., Russell, C.T., 1997. Venus: magnetic field and magnetosphere. In: Shirley, J.H., Fainbridge, R.W. (Eds.), Encyclopaedia of Planetary Sciences. Chapman and Hall, London, pp. 905–907. Moroz, V.I., Ekonomov, A.P., Golovin, Yu.M., Moshkin, B.E., San’ko, N.F., 1983. Solar radiation scattered in the Venus atmosphere—The Venera 11, 12 data. Icarus 53, 509–537. Schaefer, L., Fegley, B., 2004. Heavy metal frost on Venus. Icarus 168 (1), 215–219. Schofield, J.T., Taylor, F.W., 1982. Net global thermal emission from the Venus atmosphere. Icarus 52, 245–262. Taylor, F.W., 1981. Equatorial cloud properties on Venus from Pioneer Orbiter infrared observations. Adv. Space Res. 1, 151–154. Taylor, F.W., 1995. Carbon monoxide in the deep atmosphere of Venus. Adv. Space Res. 16 (6), 81–88. Taylor, F.W., 2006a. Climate Variability On Venus And Titan, Space Science Series of ISSI, vol. 19, Solar Variability and Planetary Climates, Springer, in press. Taylor, F.W., 2006b. On the distribution and variability of water vapour in the middle atmosphere of Venus. Adv. Space Res. in press. Taylor, F.W., Beer, R., Chahine, M.T., Diner, D.J., Elson, L.S., Haskins, R.D., McCleese, D.J., Martonchik, J.V., Reichley, P.E., Bradley, S.P., Delderfield, J., Schofield, J.T., Farmer, C.B., Froidevaux, L., Leung, J., Coffey, M.T., Gille, J.C., 1980. Structure and meteorology of the middle atmosphere of Venus: infrared remote sounding from the Pioneer Orbiter. J. Geophys. Res. 85, 7963–8006. Taylor, F.W., Crisp, D., Be´zard, B., 1997. Near-Infrared sounding of the lower atmosphere of Venus. In: Bougher, S.W., Hunten, D.M., Phillips, R.J. (Eds.), Venus 2. University of Arizona Press, Tucson, pp. 325–351. Tomasko, M., 1983. The thermal balance of the lower atmosphere of Venus. In: Hunten, D.M., Colin, L., Donahue, T.M., Moroz, V.I. (Eds.), Venus. University of Arizona Press, Tucson, pp. 604–631. Urey, H.C., 1952. The Planets. Yale University Press, New Haven. von Zahn, U., Kumar, S., Niemann, H., Prinn, R., 1983. Composition of the Venus atmosphere. In: Hunten, D.M., Colin, L., Donahue, T.M., Moroz, V.I. (Eds.), Venus. University of Arizona Press, Tucson, pp. 299–430. Williams, D.R., 2005. National Space Science Data Center. /http:// nssdc.gsfc.nasa.gov/planetary/factsheet/venusfact.htmlS.