Detection of Habitable Planets and the Search for Life

  • Heike Rauer
  • Juan Cabrera
  • Stefanie Gebauer
  • John Lee Grenfell
Part of the Cellular Origin, Life in Extreme Habitats and Astrobiology book series (COLE, volume 28)


One of the main scientific drivers for extrasolar planet research is the search for terrestrial planets in the habitable zone (HZ) and subsequently the detection of biosignatures indicating the presence of life. This goal is of fundamental importance to answer the question whether the Solar System is the only place in our universe that developed life, or if life is actually common in our galaxy and the biosphere on Earth is just one among many.


Terrestrial Planet Habitable Zone Extrasolar Planet Host Star Dwarf Star 
These keywords were added by machine and not by the authors. This process is experimental and the keywords may be updated as the learning algorithm improves.

1 Introduction

One of the main scientific drivers for extrasolar planet research is the search for terrestrial planets in the habitable zone (HZ) and subsequently the detection of biosignatures indicating the presence of life. This goal is of fundamental importance to answer the question whether the Solar System is the only place in our universe that developed life, or if life is actually common in our galaxy and the biosphere on Earth is just one among many.

However, detecting life outside the Solar System is challenging, because we are far from being able to perform any kind of in situ analysis. Current telescopes on ground and in space can provide us with detected exoplanets in the HZ of cool host stars, cooler than the Sun. Detected small, low-mass planets have to be identified as terrestrial with a rocky surface, and the presence of an atmosphere has to be established. Then, finally, atmospheric biosignatures can be searched for remotely from Earth. Techniques become even more challenging if we consider the option of habitable moons around, for example, gas giant planets.

In this section we describe the various detection methods used for exoplanets and discuss the current status of detection in particular with regard to potentially habitable planets. We then describe the methods to detect atmospheres and finally biosignatures in such planets.

2 Detection Methods of Extrasolar Planets

The detection of extrasolar planets is a challenging task since planets orbit stars which are several orders of magnitude brighter than the planet itself (Fig. 1). Due to their different temperatures, stellar and planetary emission fluxes are shifted in their peak wavelength when considering planets in the HZ. In the visible, the star is much brighter than the planet by ∼10 orders of magnitude, and in the thermal infrared (IR) spectral range, the contrast ratio of planet/stellar flux is still in the range of 10−6–10−7 for an Earth-like planet orbiting a solar-like star. For planets orbiting cool M dwarf stars with effective temperatures as low as 3,000 K, the flux ratio improves in the IR range, although the star is nevertheless much brighter than an orbiting terrestrial planet.
Figure 1.

Overview of the intensity contrast from central stars of different types (Sun, M dwarf with T eff  =  3,000 K) to an Earth-like planet with surface temperature of 288 K. For simplicity the reflected stellar light and emission fluxes are approximated as Planck curves.

Furthermore, most planets orbit too close to their host to be able to spatially resolve them from the star using current technology. Therefore, modern detection methods, with the exception of direct imaging, use indirect indicators to infer the presence of extrasolar planets. We provide here an overview of the main detection methods used. Recent textbooks describing detection techniques for further reading are, for example, The Exoplanet Handbook (Perryman, 2011) and Transiting Exoplanets (Haswell, 2010).

Radial Velocity Method

The orbital motion of the planet around its host star causes periodic wobbling of the star around their common center of mass. This wobble can be measured from the periodic Doppler shift of stellar absorption lines using stable high-resolution spectrographs. Since the orientation of the planetary orbital plane toward us is unknown, we can only detect the radial component of the motion pointing toward Earth. The method is therefore called the radial velocity method. For example, Jupiter in our Solar System induces a radial velocity (RV) signal of about 12 m/s, whereas Earth leads to an RV signal in the order of 10 cm/s only, which is below current detection limits (Mayor and Queloz, 2012). The orbital period, P p, is directly obtained from the measured radial velocity component of the stellar motion. Using Kepler’s Third law, a p 3 = GM (P p/2π)2, the planet’s orbital distance, a p, can then be determined. The shape of the RV curve, when plotting over time, provides the eccentricity, e, of the planetary orbit. The amplitude of the RV Doppler shift variation is related to the planet’s mass, m p, but also to the inclination, i, which is the angle of the projected orbital plane of the planet to the sky plane. Therefore, RV measurements provide a lower limit of the mass only, namely, the quantity: m p sin i. In the case of planetary systems with several planets, their induced RV components overlap and are disentangled during data reduction. Planetary systems with up to seven planets have already been detected so far (Lovis et al., 2011). Current detection limits of the RV method already include planets with masses in the terrestrial regime (so far mainly around stars that are lighter than the Sun). The detection of a planet with about Earth’s mass around the bright (V∼1 mag) KIV-type star α-Centauri B marks a recent milestone for RV detections. It shows also the huge effort needed for such detections, since observations lasting 4 years were needed to detect this low-mass planet on a 3-day orbit (Dumusque et al., 2012).

Here, we summarize the quantities which can be derived from the RV signal. We point out that the planetary mass is measured in relation to the stellar mass, which therefore must be derived by other means.
  • Measured quantity: Doppler shift of stellar radial velocity component

  • Derived planetary parameters: P, a, e, m p sin i

  • Required stellar parameters: m s


Another method to follow the wobbling motion of the central star induced by an orbiting planet is to follow its projected spatial motion on the sky. By using high-accuracy astrometric methods to measure the position of the host star in relation to reference stars, its periodic motion can be detected down to, for example, milli-arcseconds in the data of the Hipparcos satellite (Mignard, 1997; van Leeuwen, 2010). Unfortunately, accurate astrometric data with sufficient time sampling are sparse so far, and no planets have been discovered by this method yet. This picture is expected to change with the Gaia satellite which is anticipated to detect a large quantity of Jupiter-mass planets (Casertano et al., 2008). It needs to be pointed out that the astrometric method is particularly well suited to detect planets at intermediate orbital distances of around 1 AU. However, even Gaia’s accuracy is insufficient to detect terrestrial planets, and this detection method therefore will be limited to Jupiter-mass planets for a long time until dedicated astrometric space missions are designed (e.g., the NEAT project, Malbet et al., 2012).
  • Measured quantity: spatial motion of the star on the sky

  • Derived planetary parameters: P, a, e, m p, i

  • Required stellar parameters: m s

Transit Method

This method obtains high-precision photometric data of the stellar flux. When a planetary system is seen edge-on, the planet will periodically pass through the line of sight from Earth to its host star. During such transits, the planet occults part of the stellar disk. The stellar flux, F s, received is then reduced in proportion to the occulted area, hence in proportion to the size of the planet: ∆F α(r p/r s)2.

The amplitude of the transit signal is around 1 % flux reduction for a Jupiter-sized planet in front of a solar-like star and in the order of 0.01 % for the Earth. The transit method is the only method which allows the measurement of the radius, r p, of an extrasolar planet directly.

Since transits can be detected only for systems seen edge-on, the probability to see systems in such geometry reduces the detection efficiency of this method. For example, the geometrical probability to see a Jupiter-sized planet on a short-period orbit of a few days is around 10 % and reduces to about 0.5 % for an Earth at 1 AU. Furthermore, the duration of the transit is short in comparison to the orbital period of the planet. For example, a close-in Jupiter on a 4-day orbit transits in 3–4 h, but an Earth at 1 AU transits in 13 h only during its 1-year orbit. Thus, continuous high-precision measurements are needed not to miss such a short transit event. After detection of transit-like signals, it has to be ensured that these are not caused by other sources, for example, eclipsing binary stars and spots; hence, follow-up observations are required. The secure detection of a planet requires the determination of its mass, for example, by the RV method. If this cannot be done, for instance, if the star is too faint, at least upper limits need to be placed to separate planets from binary stars. The difficulties of the transit method are, however, more than counterbalanced by the huge potential of transiting planets for further characterization of their nature, as will be outlined below.
  • Measured quantity: photometric dimming of the stellar flux

  • Derived planetary parameters: P, a, r p

  • Required stellar parameters: r s

Timing Method

There are several ways to use timing measurements as planet indicators:
  • Transit Timing Variations (TTVs): The orbit of an observed transiting planet can be perturbed by additional planets in the system. Under certain circumstances, such perturbations will lead to variations in the time of the observed transit events. From such TTVs, it is possible to put constraints on the mass even of unseen perturbing planets based on numerical studies of the long-term gravitational stability of the system (see Fabrycky et al., 2012). In principle, even terrestrial-sized planets can produce observable perturbations (Csizmadia et al., 2010). In coplanar planetary systems where several planets show transit events, orbital periods and masses can be better constrained. An example is the Kepler 11 system with 6 transiting planets (Lissauer et al., 2011).

  • Pulsation timing: The first extrasolar planets were detected around a pulsar (Wolszczan and Frail, 1992). Their presence was inferred from periodic variations in the arrival times of pulses from the neutron star. In a similar way, periodic variations in the signal of pulsating stars can be used to infer the presence of planets. However, only the orbital parameters and an upper limit to the mass of the planet can be derived. The host pulsar is too faint to make any other characterization possible.

  • Measured quantity: central time of transit events (TTV) or pulse arrivals

  • Derived planetary parameters: P, a, e, m p sin i or upper limit on m p

  • Required stellar parameters: m s


The gravitational lensing effect (Mao and Paczynski, 1991) leads to an amplification of the stellar flux when two stars are oriented along the line of sight. If the lensing star is orbited by a planet, this again leads to an additional amplification of the background star flux, but with much shorter duration. From the properties of this amplification, the main parameters of the star and the planet acting as a lens can be derived. However, because of the relative configuration of the background and the lens star with respect to the observer, the stellar system hosting the planet is in most of the cases not observable directly (Bennett et al., 2006) and hence is not suitable for any further characterization.
  • Measured quantity: stellar flux amplification

  • Derived planetary parameters: a, P, constraints on m p

  • Required stellar parameters: (lens star is unknown, constraints in m s obtained)

Direct Imaging

Imaging of planets next to their host stars is the most direct method to detect them. However, as outlined above, it is restricted so far to planets with large orbital distances due to the limited spatial resolution of instruments and the unfavorable planet/star flux ratio (Chauvin et al., 2004; Janson et al., 2010). To improve the contrast, usually young stars are observed in the IR wavelength range, because young planets emit more strongly at these wavelengths. Models of the emitting planets then allow their mass to be derived, which is however often model dependent, in particular for very young objects. The orbital parameters are difficult to obtain because these planets typically have very wide orbits, far beyond the HZ (Chauvin et al., 2012). There are several planned projects for direct exoplanet detection (SPHERE, Beuzit et al., 2010; SPICES, Boccaletti et al., 2012).
  • Measured quantity: planet flux and spatial separation to star

  • Derived planetary parameters: m p (model dependent), in some cases P, a, e

  • Required stellar parameters: m s, distance

3 Detection Status and Future Prospects for Habitable Planet Detections

Figure 2 shows the current (April 2013) status of exoplanet detections. For comparison, the planets in our Solar System are also shown. Exoplanets discovered by the main detection techniques are indicated by different colors. Data are taken from
Figure 2.

Confirmed planet detections (March 2013). No planet candidates are included here. Red dots with downward arrows indicate planets detected by the Kepler mission, but with only upper limits on their mass available.

As outlined above, current direct imaging techniques (light purple) detect predominantly planets at large orbital distances (>2–3 AU) and are restricted to large planets in the regime of gas giants such as Jupiter. This method is therefore not yet suitable to detect planets in the HZ.

The RV measurements (light blue) provide the largest sample of planet detections so far. This technique has a long tradition in planet hunting (Walker, 2012). For distances very close to the star, RV techniques are already able to detect planets with m p sin i down to Earth mass (Pepe et al., 2011; Dumusque et al., 2012). Since these planets may be of rocky nature, they have been termed “super-Earth” up to about 10 MEarth. These close-in, hot super-Earths are a new class of objects, not found in our Solar System, similar in this respect the so-called hot Jupiters. Part of the currently known very low-mass super-Earth planets have been detected around low-mass M dwarf stars, where the more favourable mass ratio favors detections (Bonfils et al., 2013). Due to their reduced luminosity, these stars emit less flux and hence their habitable zone is closer (see Fig. 3). Thus, planets in their HZ have shorter orbital periods, again facilitating detection. For planets at 1 AU around solar-like stars, we have to await future instruments, for example, ESPRESSO (Pepe et al., 2010) and an improved data analysis treating the noise induced by stellar activity. Even with such instruments, the observational load for such detections by the RV method will be huge.
Figure 3.

Known super-Earth planets with respect to the position of the HZ which follows the scaling proportional to the solar insolation flux given by Kasting et al., 1993. Red dots with black downward arrows indicate planets detected by the Kepler mission, but with only upper limits on their mass available.

The second largest group of known planets has been detected by the transit method (red dots in Fig. 2). As expected, these detections show a clear bias toward shorter orbital distances due to the increased transit geometrical probability. However, transiting planets out to 1 AU have already been detected, demonstrating the feasibility of the method to detect planets even in the HZ of solar-like stars. More such detections are expected in future as transit searching space missions (CoRoT and Kepler, Baglin et al., 2006; Borucki et al., 2008) continue. The first detected terrestrial super-Earth with known radius and mass was CoRoT-7b (r p  =  1.6  ±  0.1  r Earth, m p  =  7.4  ±  1.2 m Earth, Hatzes et al., 2011), and the smallest confirmed planet detected by Kepler (Kepler-10b, r p  =  1.416  ±  0.03 r Earth, m p  =  4.6  ±  1.2 m Earth, Batalha et al., 2011) is even smaller and lighter. Both planets orbit solar-like central stars, although on very close-in orbits with periods <1 day.

The Kepler mission has also published a list of candidate planets (detected transits, but no RV measurements are yet available), showing that the detection of photometric transit signals corresponding to Earth-sized planets and even smaller is possible. However, confirming Earth-sized planets in the HZ of solar-like stars by RV follow-up still provides a challenge. Unfortunately, most stars observed by the Kepler mission are too faint to allow for efficient RV observations and therefore for independent confirmation of the signal and planet mass determination. Even in those cases where follow-up observations are sufficient to exclude false alarms by, for example, binary stars, only upper limits for the mass can be given if planets orbit such faint stars (e.g., red dots with downward arrows in Figs. 2 and 3). These upper limits usually do not allow one, however, to securely differentiate between a Neptune-like ice planet and a rocky, terrestrial object. Characterization of the radius and mass for such small planets around solar-like stars has to await future space missions (e.g. the PLATO mission, Catala, 2009).

Transiting planets around small M dwarfs, however, show a favorable radius ratio r p/r s, leading to larger transit depths for a given planet size. In such cases, the detection of super Earth planets comes into reach even for ground-based transit surveys. Since M dwarfs also allow for an easier detection of the RV signal, it is furthermore easier to determine their mass.

The microlensing method is most effective over a range of orbital distances from about 0.5 to 5 AU, thus covering the HZ of solar-like stars. As seen in Fig. 2, super-Earth planets have already been detected in this range (green dots). The derived planet parameters are somewhat model dependent, and detailed follow-up characterization of the detected planets and host stars is usually not possible. The main potential of the microlensing technique is to provide a large sample of low-mass planets, as well as planetary system architectures including Jupiters in relatively distant orbits for statistical planetary population considerations (Cassan et al., 2012).

The transit timing method is very interesting because it can complement the RV and transit methods, and it can potentially find very small planets in particular orbits (such as resonances). But the proper characterization of such planets requires long baselines and photometric precisions only achievable from space. At present we lack suitable candidates for terrestrial planets in the HZ discovered by this method which can be characterized with present or near-future instrumentation.

Figure 3 summarizes our current status on the detection of super-Earth planets in the habitable zone. The position of the HZ is indicated for different central star types, following Kasting et al. (1993). Known planets with m p  ≤  10 m Earth are indicated. As expected from the already discussed detection biases, most super-Earth planets are found inward of the HZ, too close to their star to be habitable. However, some super-Earths at larger distances have already been found.

The first detected potentially habitable planet is Gliese 581d. The planet was detected by RV measurements, together with three additional planets in this system (Udry et al., 2007; Mayor et al., 2009). Since this planet does not show transits, only m sin i data are available. The m sin i  =  7.09 MEarth implies that the planet may fall into the range of super-Earth planets and might be rocky. The planet is too close to its star to allow for direct imaging, and since it does not transit, spectroscopy of its atmosphere is not possible with current instrumentation. We therefore have no information on its atmospheric composition. To assess whether the planet could be habitable or not, assumptions on its atmospheric composition have to be made. Several authors have studied the potential for habitability by assuming a CO2-dominated atmosphere. This assumption is based on the atmospheres of Venus and Mars, which are CO2 rich, as well as the composition of early Earth which was probably CO2 rich for long time periods, although the detailed composition of early Earth’s atmosphere and its development with time is still debated. While the approaches of various authors differ from mere scaling to self-consistent climate modeling of the surface temperature, all conclude that Gliese 581d could be habitable for high-pressure and CO2-rich atmospheres (e.g., von Bloh et al., 2007; Selsis et al., 2007; Wordsworth et al., 2010; von Paris et al., 2010; Kaltenegger et al., 2011).

Gliese 581 is an M dwarf star (T eff  =  3,200 K). Therefore, its HZ lies relatively close to its star. Planets around such stars with more favorable viewing geometries, allowing for transits, would readily be detectable. A recent example is the detection of a super-Earth well within the HZ around GJ667C, an M dwarf which is part of a triple star system (Anglada-Escudé et al., 2012; Delfosse et al., 2012). Thus, the detection of super-Earth planets in the HZ of cool M dwarfs is certainly feasible with current techniques in the near future.

Another recent example of planets transiting in the HZ is the planet Kepler-22b. This planet was detected in the HZ of a solar-like star (Borucki et al., 2012). Unfortunately, due to the faintness of its host star, proper characterization of the mass of the planet (needed to constrain its internal structure) and its atmosphere provides a challenge for current and near-future instrumentation. With an upper mass limit of <35 m Earth and a radius of 2.4 R Earth determined from transit photometry, Kepler-22b is most likely Neptune-like rather than a terrestrial planet. Nevertheless, the detection shows that planets in the HZ of solar-like stars can indeed be achieved by the photometric transit method, despite its geometry bias toward short-period planets.

4 Detection of Habitability

Before deciding whether a detected extrasolar planet could be habitable, we have to define criteria for habitability (which is discussed in detail elsewhere in this book). Most commonly, the word “habitable” refers to surface life as we find it on Earth today, but we know from extremophiles on Earth that suitable habitable conditions can differ widely, depending on which species we are looking at. Detecting habitable conditions and life on exoplanets is challenging. The criteria for habitability of exoplanets, therefore, must be simple and robust. For these reasons, usually the presence of liquid surface water is taken as a minimum requirement. This argument is based on the assumption that life cannot develop without liquid surface water. In addition, water-based life, as we find it on Earth, is the only life form for which we reasonably know what biosignatures to expect and to search for. The criterion of liquid surface water constrains the physical surface conditions on a planet: Surface temperatures must remain above freezing for sufficiently long time periods to allow for the development of life. We emphasize that liquid surface water does not necessarily imply that life actually develops. It is merely a minimum requirement, in addition to the requirements of energy sources and nutrients.

Whether a planet can support liquid surface water depends on the heating from incoming stellar energy flux and the presence of an atmosphere including greenhouse gases. The incoming stellar radiation flux at the planet can be computed once the stellar type, hence effective temperature, and the orbital parameters of the planet are known. To constrain the surface temperature, however, requires additional knowledge about the composition and structure of the planetary atmosphere. The latter is much more challenging and beyond current detection limits for terrestrial planets in the HZ of Sun-like stars, as further outlined below.

The orbital distance of a planet to its host star is used as a first indicator for potential habitability. Kasting et al. (1993, Fig. 3) determined the orbital limits of the HZ for Earth-like planets. The orbital position of the HZ scales mainly with the incoming stellar energy flux and the HZ moves closer to the host for cooler stars. For a particular terrestrial extrasolar planet, its orbital distance can be compared to these HZ limits. If the planet is within these limits, it is often called potentially habitable. However, this implies the assumption of a terrestrial planet with a suitable atmosphere. As an example, we note that the Earth’s moon is also in the HZ of the Sun, but is certainly not habitable.

To better constrain the potential habitability of planets, atmospheric models are used. For this approach assumptions on the atmosphere of the terrestrial planets have to be made. Often, Earth-like atmospheres are assumed, but also CO2-rich atmospheres are considered (e.g., Segura et al., 2003, 2005, 2010; Grenfell et al., 2007a, b; Kitzmann et al., 2010; Rauer et al., 2011). Should atmospheric compositions of terrestrial exoplanets be detected in future, such model scenarios obviously can be much better constrained and predictions for particular exoplanets be made. Such predictions form the basis of the search for biosignatures on such planets, but also for the interpretation of measured atmospheric signals.

The detection of an ocean glint, the forward reflected stellar flux on an ocean, has been proposed as a way to detect surface water on exoplanets directly (Williams and Gaidos, 2008; Robinson et al., 2010; Barry and Deming, 2011). However, again this detection will be extremely challenging and difficult to disentangle from other planetary features and stellar activity.

In summary, the main criteria for the potential habitability of an extrasolar planet are its orbital distance, the central star type, determining the stellar energy input to the planet, and the composition and structure of its atmosphere, determining the greenhouse heating effect at the planetary surface.

5 Detection of Atmospheres and Biosignatures

A biosphere on extrasolar planets can only be detected if it produces a sufficiently large signal in the atmosphere. Due to the large distance to Earth, only planetary disk-integrated fluxes can be observed on exoplanets. A further complication arises from the nearby host star whose emission flux has to be separated from the faint planetary signal. Extrasolar planets orbiting at a sufficiently large distance to their host star can be detected through high-spatial resolution imaging methods, and their emitted infrared (IR) flux can be observed directly. Unfortunately, so far only large gas giants orbiting at distances outside the habitable zone can be detected by this method (see Fig. 2). Plans in Europe and US have been made over the past decades to obtain images and spectra from nearby terrestrial exoplanets, for example, with the Darwin (Leger et al., 1995) and TPF (Beichman et al., 1999) missions. Both projects were identified as too complex to achieve at present, although similar missions are likely to be re-proposed in the future as the need to search for biosignatures increases with the increasing number of terrestrial exoplanet detections.

The detection of atmospheres is so far limited to planets transiting in front or behind their host star (Fig. 4), or via the variation of the reflected stellar light with orbital phase of the planet. When the planet transits in front of its star, the stellar flux passes through the optically thin parts of its atmosphere. When comparing spectral measurements of the (planet  +  star) flux to those times when the planet is not in the line of sight, the stellar spectrum can be separated from signatures of the planet atmosphere. These measurements are usually performed in the optical wavelength range where the stellar flux is largest and high signal-to-noise ratios (SNRs) can be obtained. The planetary atmosphere can also be detected by its IR emission flux. For this purpose, the (planet  +  star) flux is observed just before the planet vanishes behind the star during secondary eclipse (Fig. 4). Comparing to measurements during the planetary eclipse (obtaining only the stellar flux) again allows us to separate the stellar and planetary flux. Both methods have already successfully been applied to gas giant and Neptune-sized planets (e.g., Deming et al., 2005; Richardson et al., 2007; Tinetti et al., 2007; Knutson et al., 2008; Pont et al., 2008; Redfield et al., 2008; Swain et al., 2008; Bean et al., 2010; Beaulieu et al., 2011; Croll et al., 2011; Crossfield et al., 2011; de Mooij et al., 2012).
Figure 4.

Geometry of transiting extrasolar planets. The atmosphere of the planet can be probed during primary transit, when the stellar light passes through the optically thin parts of the atmosphere, and during secondary eclipse, when the emitted infrared flux of the planet can be detected. During the orbit, reflected stellar light varying according to the orbital phase of the planet can be observed.

Chapter  5 discusses which biomarker signals in planetary atmospheres can be interpreted as indicators for habitability and life. The presence of water (a mandatory ingredient for life), carbon dioxide, oxygen, and ozone is taken as indicator for life on our planet. However, we also have to be aware of false positives or negatives, e.g. via abiotic sources mimicking the signatures of life, or masking the presence of biomarker signals which then remain undetected in exoplanets. Since the interpretation of atmospheric biosignatures is therefore challenging, most studies to simulate the spectral appearance of exoplanets assume an Earth-like biosphere as a first approximation whose parameters we know best (e.g., Selsis et al., 2002; Segura et al., 2003; Kaltenegger et al., 2007, 2011; Kaltenegger and Traub, 2009; Kaltenegger and Sasselov, 2010; Kitzmann et al., 2011a, b; Rauer et al., 2011).

The detection of biosignatures in atmospheres of terrestrial planets orbiting in the HZ around solar-like stars is a highly important goal of exoplanet research, aiming at a better understanding of our Solar System in comparison to other similar systems. However, in view of the diversity of planetary systems found and because of the more favorable contrast ratio between planet and host star (Fig. 1), terrestrial planets orbiting cool M dwarfs have gained interest in recent years. Furthermore, the probability of detecting transiting planets in the habitable zone is larger for planets orbiting M dwarfs, because the HZ is closer to the star and the star is smaller compared with the Sun. This further strengthens the interest in M dwarfs. Figure 5 shows as an example the modeled transmission spectrum of the Earth passing in front of the Sun during primary transit (black line). The typical atmospheric spectral absorption bands of several molecules in the Earth’s atmosphere are seen: for example, H2O, CO2 and CO, CH4, N2O, and O3. The figure also illustrates how the spectral appearance of the same planet changes when it orbits in the HZ of different M dwarf host stars. M dwarfs are small, cool stars which emit their energy spectrum more to the red spectral range (maximum flux around 1–2 μm, Fig. 1) when compared to the Sun. Quiet M dwarfs have very low UV emission fluxes, in particular at wavelengths where atmospheric chemical processes are relevant. However, active M dwarfs, like AD Leo in Fig. 5, can also emit the same or even more UV spectral flux as the Sun due to their strong activity. The different spectral energy flux distribution of M dwarfs has an effect on the chemical processes in the atmosphere of the orbiting planet, even though the total stellar energy received by the planets from the different host stars is the same for all examples shown. As a result, a planet with an atmosphere like modern Earth placed at an orbital distance around an M dwarf where it receives the same total energy as Earth can appear spectroscopically very different. This can even mean that spectral signatures of the biosphere disappear, leading to false negatives (Rauer et al., 2011). This example shows that interpretation of observed spectra of terrestrial planets has to be made with care.
Figure 5.

The “relative transmission,” hence the additional absorption during transit due to the planetary atmosphere, is shown. The modeled atmospheric transmission spectrum of the modern Earth passing in front of the Sun is shown in black. In addition, models of the same planet around the active M dwarf AD Leo (red) and two quiet M dwarfs of type M5 (blue) and M7 (magenta) are shown (From Rauer et al., 2011).

To estimate whether the detection of such atmospheric absorption bands on terrestrial super-Earth planets is actually feasible, one has to consider the stellar and planetary flux received on Earth and the instrument used for detection. In particular in the near IR, where many interesting atmospheric absorption bands appear, the intensity contrast ratio between planet and star is relatively high. In addition, beyond about 2 μm intensities drop significantly, again forming a challenge for instruments.

Biosignatures have also been suggested when observing light reflected by the planet in the visible wavelengths range (Fig. 1) where, for example, the signature of the red-edge indicating chlorophyll on Earth has been proposed as biosignature also for exoplanets (e.g., Arnold et al., 2009; Briot, 2010). Some studies suggest using polarization to improve the contrast between star and planet (e.g., Seager et al., 2000; Saar and Seager, 2003; Stam, 2008), but again the low flux levels represent a challenge.

Several SNR estimates have been made in the literature for instruments on the future James Web Space Telescope (JWST) (e.g., Kaltenegger and Traub, 2009; Shabram et al., 2011; Belu et al., 2011; Rauer et al., 2011). A straightforward approach was taken, for example, by Belu et al. (2011) using Planck functions to approximate the stellar and planetary flux and assuming reasonable absorption band strengths derived from consistent atmospheric model calculations (as, e.g., in Fig. 5). This work considered transmission and emission spectra during primary transit and secondary eclipse observed with two instruments on JWST, namely, NIRSpec and MIRI. Noise introduced by these instruments and from the zodiacal light background of the sky (scattered light on interplanetary dust particles in our Solar System) was considered. Model simulations included a wide range of central stars, from Sun-like to M dwarfs. It was shown by this study that transiting habitable super-Earths can be characterized at low spectral resolution with a significant SNR only in the most nearby systems (<10 pc) and when hosted by a cool M dwarf star (<0.2 Msun) (e.g., Fig. 6 for SNR of the 9.6 μm ozone band). Similar results were also obtained by Rauer et al. (2011) and Kaltenegger and Traub (2009), again showing that many transit observations of terrestrial planets in the HZ have to be co-added to reach a reasonable SNR, even with the 6.5 m mirror JWST telescope. These modeling studies provide an example. Actual absorption strengths of atmospheric bands depend on mixing ratios, temperature profiles, incoming stellar flux, and chemistry. However, all studies show that the detection of atmospheric absorption bands in terrestrial exoplanets is indeed challenging, already for the main atmospheric constituents, and even more so for less abundant biosignatures.
Figure 6.

Modeled signal-to-noise ratio (SNR) for detection of the 9.6 μm ozone band for a terrestrial planet orbiting M dwarf host stars located at 6.7 pc from Earth (From Belu et al., 2011). SNRs are given for transmission observations during transit with the MIRI instrument at the JWST. All observable transits during the planned 5-year mission lifetime of JWST have been co-added to obtain the SNR shown. The gray band indicates the position of the HZ. Top, primary transit; bottom, secondary transit.

Further complications for the detection of atmospheres arise when clouds are present in the atmosphere. The recently detected mini-Neptune-like planet GJ1214b orbiting an M dwarf star (Charbonneau et al., 2009) can serve as an example here since it is large enough to actually allow for atmospheric signature detections with current instruments. However, all data points obtained over a wide wavelengths range from visible to IR show a featureless spectrum which is interpreted as cloud cover in the planet’s atmosphere masking spectral absorption bands that could otherwise indicate the composition of the atmosphere (Bean et al., 2010). Similar problems can occur for terrestrial planets with significant cloud cover (Kaltenegger et al., 2007; Kitzmann et al., 2011a).

Planetary systems around M dwarfs are nevertheless very interesting to investigate, since the habitable zone is close to the star (Fig. 3) hence planets there can be more easily detected, as already outlined. Furthermore, the population of stars in our neighborhood is in general dominated by M stars. However, model simulations of expected SNRs have shown, as discussed above, that only the nearest M dwarfs (within 10–20 pc) are feasible for detections of spectroscopic atmospheric signals. Only a small number (∼300) of M dwarf stars can actually be found within 10 pc from Earth, limiting the potential number of target stars. Based on a recent RV survey of M dwarfs (Bonfils et al., 2013), about 100 of them are expected to have planets orbiting in the HZ. With a geometrical probability of 1–3 % for transits of such planets (Kaltenegger and Traub, 2009), only 1–3 planets transiting around nearby M dwarfs within 10 pc are expected to be detected in future. Therefore, the number of detected transiting terrestrial planets suitable for atmospheric characterization around such nearby M dwarf host stars will unfortunately remain limited.

Future space mission plans include dedicated telescopes for exoplanet atmosphere detections during planet transits, for example, EChO (Tinetti et al., 2012). Such missions concentrate on Jupiter- and Neptune-like planets and focus for terrestrial planets on M dwarf host stars to take advantage of the improved flux ratio between star and planet and the close-in HZ. However, due to the limited number of very nearby stars, the potential target number for terrestrial planets around these stars will be small.

Finally, we point out that stellar activity and related energetic radiation fluxes of M dwarfs may be a challenge for planets in their close-in habitable zones. It is by no means clear that terrestrial planets around cool M dwarfs exhibit an Earth-like biosphere. The next step along the detection of biospheres therefore is clearly to target Solar System analogues where interpretation of atmospheric signals in terms of biosignatures will be facilitated, and furthermore allow us to put our system into the context of similar systems in our galaxy. But this will require space-based, direct imaging telescopes that are currently not within projected mission plans for either ESA or NASA.

6 Summary

The detection of habitable planets outside our Solar System is a challenging goal. However, recent years have shown that it is feasible to find terrestrial super-Earth planets in and near the habitable zone of M dwarfs and even solar-like stars in the near future. This search is facilitated for cool M dwarfs because their HZ is close to the star which makes the detection of planets easier, for radial velocity detections as well as for planetary transits. Furthermore, the improved planet/star flux contrast in principle helps the detection of atmospheric signatures. On the other hand, model simulations have shown that biosignatures of Earth-like planets can only be detected for planets transiting very nearby M dwarfs (<10 pc), which are unfortunately sparse. In addition, it is by no means clear whether active M dwarfs would allow for the formation of a biosphere on planets in its HZ. Searching for planets more similar to Earth requires that we first detect terrestrial planets in the HZ of solar-like stars. Such detections are in principle feasible for near-future radial velocity instrumentation, albeit extremely time consuming. As a result, their numbers will remain small. Future transit surveys targeting terrestrial planets in the HZ of bright solar-like stars, like the PLATO mission, have the potential to significantly increase these statistics. Other methods targeting planets in the HZ in future are direct imaging, astrometry and microlensing. Detecting biosignatures on any suitable terrestrial planets found is an even more challenging goal. With current instruments, this seems feasible only for the brightest and nearest stars. To obtain a sufficiently high SNR, such instrumentation will need a large aperture telescope with the means of separating the planet from the star, for example, by coronographs, starshades, occulters, or interferometry, and forms one of the main challenges for future exoplanet research.


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Copyright information

© Springer Science+Business Media Dordrecht 2013

Authors and Affiliations

  • Heike Rauer
    • 1
    • 2
  • Juan Cabrera
    • 1
  • Stefanie Gebauer
    • 2
  • John Lee Grenfell
    • 3
  1. 1.Planetary ResearchGerman Aerospace Center (DLR), Institute of Planetary ResearchBerlinGermany
  2. 2.Zentrum für Astronomy und AstrophysikTechnische Universität Berlin (TUB)BerlinGermany
  3. 3.Institut für PlanetenforschungDeutsches Zentrum für Luft- und Raumfahrt (DLR)BerlinGermany

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