Fig. 7.22 Abundances of rare gases and their isotopes in the atmospheres of the terrestrial planets and in meteorites. A solar composition would be represented here by a straight, horizontal line. Except for xenon, the rare-gas composition of Earth and Mars resembles that of carbonaceous chondrites. That of Venus is closer to the solar wind implanted in meteorites (but it should be noted that the measurements of the atmosphere of Venus are very uncertain). This diagram shows that the atmospheres of the terrestrial planets are the result of degassing of the volatile component of solid materials accreted during their formation. This volatile component is depleted and fractionated by escape processes (Zahnle et al., 2007)
of the terrestrial planets is therefore, by its very nature, extremely different from the gas in the protosolar nebula, and the composition of rare gases is similar (in relative abundances) to that of carbonaceous meteorites (see Fig. 7.22). So starting with cosmic abundances of elements, as is done to model the giant planets, does not make any sense when it comes to the terrestrial planets. One possible approach would be to take the mass abundances of volatile components (H, C, O, and N) trapped in meteorites and from those deduce the mass and composition of elements in the atmosphere. This would, however, neglect the escape phenomena that took place during the formation and primordial evolution of the atmosphere, as well as the division of the gases between the planetary interior and the atmosphere. The Earth's interior (mantle and crust) contains approximately 5 x 105 times more CO2 than the atmosphere, most being held within the mantle or stored in sedimentary rocks in the form of carbonates. In the case of Venus, a significant fraction, if not almost all of the CO2 is found in the atmosphere (90 bars) and corresponds more or less to the whole of the terrestrial reservoir, estimated at about 150 bars. This difference in distribution is explained by the absence of liquid water on Venus, because this allows the formation of carbonates at low CO2 pressures, and also the functioning of the terrestrial carbon cycle.
In addition, within a single protoplanetary system, the planetesimals taking part in the formation of the terrestrial planets do not all have the same content of volatile components. The abundances of water, carbon, and nitrogen increase significantly with orbital distance. In the Solar System, meteorites originating in asteroids orbiting at less than 2.6 AU are water- and gas-poor. Those originating in more distant zones contain up to 15 per cent of water in the form of hydrated minerals. Bodies formed outside the ice line (4-5 AU, see Sect. 188.8.131.52) consist, on average, of 50 per cent water ice. From which of these populations of planetesimals did Earth's atmosphere originate? This question is still subject to considerable debate, but it seems that the majority of the population was poor in volatiles (ordinary chondrites or enstatites), and the water and gas primarily came from a lesser contribution of carbonaceous-chondritic material originating in the outer region of the asteroid belt (2.6-3.5 AU) (Raymond et al., 2007 and Sect. 184.108.40.206). The cometary contribution (ice) is thought to be very minor for the Earth. So it seems that the existence of oceans and a dense atmosphere on Earth depends in a crucial manner on a small contribution by mass of planetesimals formed relatively far away from the Earth's orbit. This contribution of material rich in volatiles during the accretion phase is likely to vary quantitatively from one terrestrial planet to another inside a single planetary system, and also from one planetary system to another. This contribution is indeed particularly sensitive to the location and properties of the giant planets, which formed before the terrestrial planets; the diversity of which has already been shown by recently detected exoplanets. To sum up, two terrestrial planets of the same mass, forming at the same orbital distance from two stars of the same type, could well draw on different reservoirs (in mass and composition) for their atmospheric components. How, under such conditions, can we hope to predict and model the atmospheres of terrestrial-type exoplanets? We shall see that understanding of the Earth's geochemical cycles have allowed us to define the concept of a habitable planet. Theoretical studies of the properties of habitable planets offer a point of departure for exploring the diverse types of terrestrial planets that may be expected, and also provides a fascinating link between exoplanets and astrobiology.
A habitable planet is taken to be a terrestrial planet whose surface is covered, at least partially, by a stable layer of liquid water, over geological periods of time. The habitable zone around a star corresponds to the region around a star where such planets may exist. The presence of liquid water is a condition judged to be necessary for the existence of life, but which is probably not sufficient. This statement may, however, be called into question if it were to be shown that life can develop in the absence of liquid water, by making use of other fluids. This would mean a completely different biology from the one that we know, so for operational reasons, for the moment we will limit our considerations to life as we know it on Earth. Certain locations in the Solar System, other than Earth, undoubtedly conceal non-surface liquid water, such as the subsoils of Mars and the icy satellites (Europa and perhaps Callisto), or clouds on Venus and nothing a priori forbids forms of life from existing there. However, we are interested here in searching for signs of life on planets around other stars. The first space observatories that may allow us to detect habitable worlds around other stars, such as the Darwin and TPF projects, will have a chance of detecting signs of biological activity on a planet only if the biosphere is sufficiently widespread and active to alter the observable factors (the overall composition of the atmosphere or the surface). The presence of liquid water at the surface allows the existence, in particular, of photosynthetic activity (liquid water and starlight being available simultaneously). The possibility of using stellar photons and water as an electron donor allows the biosphere to transform the planetary environment in a way that is quantitatively and qualitatively distinct from other metabolisms (Rosing 2005). A planet harbouring a restricted biosphere that is relatively inactive or confined beneath the planetary surface without modifying the surface characteristics, will never reveal the presence of life to observation from a distance. In such a case, only exploration in situ would detect biological activity. Future exobiological missions for exploring Mars and Europa, will be directed towards this end. For this reason, an exoplanet that lacks the possibility of having liquid water at its surface would be considered 'uninhabitable', even though liquid water and life forms might actually exist in its interior.
a. The circumstellar habitable zone
The habitable zone around a planet is defined as the region in which water can exist permanently in the liquid state at the surface of the planet. This implies a surface temperature Ts above 273 K and sufficient atmospheric pressure. Ts depends on the incident stellar energy and its spectral distribution, the radiative properties of the atmosphere (absorption, scattering, and emission) and the reflective properties (the albedo) of the surface and of the atmosphere (primarily through the existence of clouds). Kasting (1988), and then Kasting et al. (1993) have applied realistic atmospheric models and estimated the limits of this zone for different types of star. These authors have decided to tackle this complex problem by considering the Sun-Earth model: by modifying the orbital distance they have been able to determine the minimum and maximum distances at which the Earth's oceans are, respectively, completely evaporated or completely frozen.
b. The inner limit of the habitable zone
Currently, most of Earth's surface water is contained in the oceans. Decreasing the Earth's orbital distance consequently vaporizes a greater fraction of this reservoir of water. The increase in the level of water vapour in the atmosphere produces a more effective greenhouse effect and thus creates additional heating of the surface, provoking the evaporation of yet more water. Happily, more water vapour in the atmosphere also implies more effective cooling of the surface by convection. This second effect is essential, because, without it, the greenhouse effect on our planet would run away, and would lead to all the oceans boiling away. Taking these two mechanisms (greenhouse effect + convection) into account in what is known as a radiative-convective model allows us to determine the temperature and water-vapour profile of the atmosphere.
From surface temperatures of about 370 K, corresponding to a distance of 0.95 AU (from the current Sun), water vapour becomes the dominant component in the atmosphere. If the planet is brought even closer to the Sun, the amount of water vaporized becomes very considerable and the energy emitted by the planet in the infrared reaches a plateau, because, thanks to the high opacity of an atmosphere of H2O, the zone in which the thermal emission arises is displaced towards a radiatively isothermal region at a high altitude and becomes unaffected by the surface temperature. The thermal energy radiated away at Ts = 550 K or at Ts = 1500 K is almost the same. It is only when Ts > 1500 K that the thermal emission climbs once more, together with Ts, because the surface then radiates in the near infrared, and then cools directly to space through atmospheric windows that are transparent at such wavelengths. In terms of planetary evolution, the threshold corresponds to a dramatic transition: if a planet reaches Ts = 550 K, any increase, however minor, in the stellar luminosity (or decrease in the orbital distance) causes heating of the surface to 1000 K, and the complete evaporation of the reservoir of water. This is probably the fate that Venus has suffered early in its existence. This phenomenon, known as the runaway greenhouse effect is illustrated in Fig. 7.23.
The orbital distance of 0.84 AU calculated by Kasting should be considered as indicative only, because this estimate has been obtained for a cloud-free atmosphere, and also because there are uncertainties about the spectroscopic properties of water vapour at high pressure. The presence of clouds seems realistic for an atmosphere dominated by water vapour, and would limit the intensity of heating of the surface. Indeed, clouds effectively reflect solar radiation: Venus, for example, reflects 70 per cent of the energy received because of its global cloud blanket, whereas without clouds, the albedo would not exceed 30 per cent. An albedo of 70 per cent would allow a liquid ocean to be retained out to orbital distances of about 0.5 AU.
Moreover, for orbital distances less than 0.95 AU, water vapour becomes the main component of the atmosphere at all altitudes, except in the very highest layer, where H2O is dissociated into H and O by UV radiation. In Earth's present-day atmosphere, water vapour is an extremely minor component above the tropopause. The tropopause is the boundary between the troposphere and the stratosphere, where the temperature reaches a minimum (see Sect. 220.127.116.11). It forms a 'cold trap' that water vapour cannot cross, where it condenses as it rises. This cold trap keeps the loss of hydrogen to space to negligible values, this being limited by the molecular diffusion of H towards the exobase (the altitude at which collisions become negligible). In a hot atmosphere dominated by H2O, photodissociation of H2O and the escape of H are no longer limited other than by the EUV irradiation of the upper atmosphere. Under these conditions, an amount of water equivalent to the terrestrial ocean may be lost in less than 109 years. It is estimated that this is the fate that Venus has undergone. Early in its history, the solar luminosity was sufficiently weak for a very hot, liquid ocean to exist, as shown in Fig. 7.23. However, the erosion of this reservoir of water by UV radiation must have been very effective, and if Venus had a quantity of water less than that in Earth's oceans, the latter must have been lost even before the increase in solar radiation reached the threshold at which total vaporization occurred. The quantity of water present on a planet therefore needs to
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