How May Life on an Exoplanet be Detected

Given the distance that separates the Earth from even the closest stars (apart from the Sun, the star closest to the Earth, Proxima Centauri, is about 4.3 light-years away), it seem inconceivable that one could contemplate searching for life, in situ, at any time in the near future. Any such 'on site' exploration assumes:

• identification of a target where the presence of life has almost been established;

• mastery of sending high-speed interstellar missions - at one tenth of the speed of light, which would mean a real feat of interstellar navigation - it would take about 45 years to reach Proxima Centauri.

as a result, the lead time for such a mission would be several generations, with the risk that any probe would immediately become obsolete right after its launch or at least during its theoretical lifetime.

To detect life on an exoplanet, we will therefore consider only those methods that involve doing so from at distance, by observation of the object. That therefore supposes that life on a planet should be effectively detectable at a distance, and that it should have developed on a planetary scale. In particular, if life has developed in a very restricted ecological niche, such as purely locally, at the bottom of a lake, for example, or in a specific environment (with specific requirements for acidity, salinity or any other particular condition), it could well be un-

detectable. We will also assume that life as it developed on a planetary scale has modified the planet's environment, such as the composition of the atmosphere, or the electromagnetic environment. The latter depends on whether life has evolved into a form that is capable of communicating through electromagnetic radiation. (see Sect. 9.7.3).

Here, we shall deal solely with the question of detecting life through a search for spectral signatures (subsequently called 'biomarkers' or biological signatures) in the planet's atmosphere.

Such spectral signatures are obtained by spectroscopy from the light coming from the planet. This technique generally requires direct observation of the planet (separating the photons from the planet from those emitted by the star), and several projects have been described that would be capable of doing so (DARWIN, TPF, etc.). However, the differential spectroscopic analysis of the light from the star before and during a planetary transit - an indirect method of detection - has already allowed the detection of several species in the atmosphere of HD 20945B, thanks to the HST, in particular Na, H, C, and O (Charbonneau et al., 2002; Vidal-Madjar et al., 2004). The method may also be applied to secondary transits (behind the star), or through observations made in the thermal infrared in phase with the object's period of revolution. In this last case, the infrared excess is measured before the eclipse, relative to the star's flux during the eclipse. Deming et al. (2005) and Charbonneau et al. (2005) have used this technique with the Spitzer space telescope to carry out photometry in four spectral bands of HD 20945B and Tr-ES 1. This technique, if carried out by the JWST, should allow us to just about obtain low-resolution spectra of some giant planets that transit. Given the low signal-to-noise ratio of the method (because of the significant infrared flux from the star), it is limited to the observation of giant planets, and spectroscopy of terrestrial-type planets requires direct observation. Spectral Signatures of Life

Any choice of biological signatures assumes a preliminary definition of what called life, and the theories on which the definition is based (see Sect. 9.1). This subject on its own could fill a complete book.

All the classic definitions of life cite a structured, material system, possessing a certain coded content of information. Chemical coding is one of the most obvious solutions, in particular, one based on the chemistry of carbon. The main reason is that on Earth, but also in the interstellar medium, carbon exists abundantly in several oxidized (e.g., CO2) and reduced (e.g., CH4), forms. This allows the formation of a considerable number of different combinations (molecules), which in turn simultaneously allow very diverse forms of coding, and of encoded structures. In addition, it is relatively easy to pass from one to the other by processes that require a limited amount of energy.

If we assume life where the coded information is based on carbon chemistry, then a large quantity of carbon in reduced form (i.e., organic molecules) is required.

However, the principal basic material is CO2, that is, an oxidized form. The reaction required to turn oxidized carbon into an organic carbon compound is carried out on Earth by cyanobacteria and plants, and is photosynthesis, as mentioned earlier in this section. This reaction results in the release of oxygen.

Because the reaction is reversible (organic matter oxidizes), and because at the same time oxygen can oxidize other, highly reducing species (such as volcanic gases, metals, rocks, etc.) the massive presence of oxygen in an atmosphere may be considered as proof of the existence of biological activity. This point was put forward by Owen in 1980, and used in 1993 by Sagan and his colleagues when the Galileo probe flew past the Earth, to detect life on Earth from a distance by observing the simultaneous presence of oxygen (the oxidant) and methane (the reducing agent) - two species that could not exist in equilibrium - in the Earth's atmosphere. At the same time, if all biological production of oxygen were to cease, it would take only about 5 to 10 million years for the atmospheric oxygen to be consumed by the oxidation process and for it to disappear from the air.

Since 1980, several studies have been carried out to evaluate this criterion, by attempting to see if it is possible to generate, by abiotic means, large quantities of oxygen that could persist in the atmosphere of a planet. In particular, photolysis of CO2 has been suggested (Rosenqvist and Chassefiere, 1995), as well as the photolysis of water, through a runaway greenhouse effect. Selsis (2008) has shown, in particular by studying the sensitivity of chemical species to their environment (type of star, distance from the star, etc.) that the criterion should be modified and that to obtain a reliable criterion, rather than just the detection of O2, the triple detection of CO2, H2O and O2, or even better, 03, would be preferable as a marker for oxygen. Oxygen (O2), in fact, has no spectral signature in the thermal infrared. Ozone (O3), is a dissociation product of O2 caused by ultraviolet photons from the star (via the Chapman cycle). Ozone is a good marker of the quantity of oxygen in an atmosphere. (There is a complete description of the discussion of the relevance of O3 as a marker for O2 in Leger et al., 1993, and in Selsis et al., 2008.)

Several other biological signatures have been suggested, but do not appear to be as effective as the triple detection of CO2, H2O, and O3:

• As mentioned earlier, the simultaneous detection of species out of thermody-namic equilibrium may be proof of biological-type activity, for example CH4 and O2, or CH4 and H2O. It is, however, difficult to extract a simple, applicable, spectroscopic criterion, without an in-depth knowledge of the planet's atmosphere and without modelling the different interactions.

• The detection of technological gases (gases the synthesis of which is impossible or highly unlikely naturally), seems to be a good tracer of biological activity, but requires evolution of a technological civilization. Study of such tracers therefore resembles a SETI-type search (see Sect. 9.7.3).

The question of the composition of planetary atmospheres and the corresponding spectra has been treated in more detail in Sect. 7.3. Choice of Spectral Region

The choice of the spectral region for observation is a compromise between:

• the presence or absence of spectral signatures of the elements being sought in the region concerned,

• the ease with which these spectral signatures may be detected (width of lines, planetary flux, and stellar flux).

The contrast between the electromagnetic spectrum of Earth and the Sun has already been mentioned in Chap. 2. It may be shown that this contrast is minimal in the radio region, but in that case the planetary flux is extremely low.

The planetary flux received at the Earth is a maximum in the thermal infrared (about 10 photons per second per square metre of detector for a planet like the Earth at a distance of 10 parsecs, observing in the 6-20 |m spectral band). In such a case, the contrast with the star is about 7 million. The thermal infrared is also a region suitable for the detection of all the other atmospheric gases, because it is the wavelength region that corresponds to the vibration of the molecules. So it is possible to envisage the detection of CH4, NO, NO2, SO2, and all the asymmetrical molecules that are present in large quantities in the atmosphere.

It is also possible to envisage being able to detect the three species CO2, H2O and O2 in the visible/near infrared region. Here, the number of chemical species that may be identified is reduced (being the band in which there are the harmonics of the fundamental frequencies of the various molecules), the planetary flux is weaker (by a factor of about 30 for the 0.5-2 |m band relative to the 6-20 |m band), and the contrast is higher by a factor of about 1000. However, certain projects with dedicated coronagraphs that are adapted to large-size telescopes (of a few metres) are capable of spectroscopic observation of terrestrial-type planets in the visible region.

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