Habitable Planets: Targets and their Environments

**FULL TITLE** ASP Conference Series, Vol. **VOLUME**, **YEAR OF PUBLICATION** **NAMES OF EDITORS** Habitable Planets: Targets and their Environments M. Güdel ETH Zurich, Institute of Astronomy, 8093 Zurich, Switzerland Abstract. The environment in which a planet resides is largely controlled by the host star, not least by its magnetic activity. The resulting interplay between the environment and a planet may include magnetic interactions, interactions between stellar radiation (optical, UV, X-rays) and planetary atmospheres, and interactions between stellar particle fluxes (high-energy particles, stellar winds) and planetary magnetospheres. 1. Introduction Habitability of planets is most significantly determined by the radiation environment defined by the host star. Kasting et al. (1993) defines habitability based on the criterion of the potential presence of liquid water on a planet bearing an Earth-like atmosphere; the inner edge of the habitable zone (HZ) is defined by loss of water as a consequence of photolysis and hydrogen escape, and the outer edge is determined by the formation of CO 2 clouds cooling the planet by the increased albedo, leading to runaway glaciation. The exact location of the HZ of course depends on the atmospheric composition, in particular the presence of greenhouse gases such as CO 2 or CH 4. Further conditions contribute to the potential habitability or inhospitality of planets. Specifically, the presence of plate tectonics has been recognized to be of prime importance for climate stabilization through the carbonate-silicate cycle (Walker et al. 1981; Kasting & Catling 2003). Also, habitable zones vary on evolutionary time scales of the host star, the slow increase of a main-sequence star s luminosity leading to increasing radii of the HZ. For example, the zeroage main-sequence Sun s bolometric luminosity was about 30% lower than the present-day Sun s (Sackmann & Boothroyd 2003), implying less insolation of the young planetary atmospheres in our solar system (Sagan & Mullen 1972). However, actual habitability conditions on planets also depend on a number of further properties of the host star not related to the relatively stable optical output of the solar surface. Planetary atmospheres are particularly susceptible to short-wavelength/high-energy radiation from the outer layers of the host stars, induced by magnetic activity. Further, late-type main-sequence stars are assumed to continuously lose mass in ionized, magnetized stellar winds that strongly interact with planetary atmospheres or magnetospheres. In contrast to the optical output from a main-sequence star, the short-wavelength radiation is subject to very strong variations on evolutionary time scales, on time scales related to so-called activity cycles, on daily to monthly time scales reflecting the changes in the magnetic configuration on the star, and on hourly time scales 1

2 related to flares. These variations in turn are a function of stellar mass. The combination of the optical luminosity (determining the classical HZ for the potential presence of liquid water) and the magnetic activity behavior of the star (driven by an internal dynamo and being a strong function of stellar age) leads to a number of important constraints for habitable planets. The following sections summarize issues related to stellar magnetic activity relevant to the habitability of planets, and discuss implications for the conditions for life. 2. The Radiation Environment in Stellar Evolution The Sun has significantly brightened during its evolution on the main sequence, starting on the zero-age main sequence at a level of approximately 70% of its present-day bolometric luminosity (Sackmann & Boothroyd 2003). The lower luminosity of the young Sun poses a problem for our understanding of the evolution of young planetary atmospheres as both Earth and Mars should have been in deep freeze during the first few 100 Myr (Sagan & Mullen 1972; Kasting & Catling 2003). In contrast, geological records indicate a warmer climate and the presence of liquid water on both planets at ages younger than 1 Gyr. Although greenhouse gases may significantly elevate atmospheric temperatures, even a high-pressure CO 2 rich atmosphere on Mars would not have produced a sufficiently mild climate (Kasting 1991; Sackmann & Boothroyd 2003). However, upper atmospheres of planets are much more susceptible to changes in the short wavelength (ultraviolet [UV], far-ultraviolet [FUV], extreme-ultraviolet [EUV], and X-ray) radiation due to stellar magnetic activity. Opposite to the bolometric luminosity of a star, its magnetically induced high-energy luminosity decreases with time as the internal dynamo decays as a consequence of stellar spin-down by angular momentum loss. The short-wavelength decay is much more dramatic than the optical increase, the X-ray luminosity decreasing over three orders of magnitude during the main-sequence life of a star; the radiation also becomes much softer on evolutionary time scales (e.g., Güdel et al. 1997). The decay time scale of magnetic activity is a function of stellar mass. While the X-ray luminosity of a a solar-mass star drops by a factor of 100 during the first Gyr, stars of spectral type M may reside at a high, saturated magnetic activity level during this entire period, after which a rather slow decay sets in (West et al. 2008). The decay laws are also a function of wavelength; shorter wavelength radiation decays by larger factors than longer wavelength radiation. Ribas et al. (2005) studied the long-term evolution of radiation shortward of 1700 Å to construct a comprehensive model of the spectral evolution of the Sun in Time. The integrated irradiances correlate tightly with the stellar rotation period or age, the relations being represented by power laws, as illustrated in Figure 1. From band-integrated luminosities and individual lines (see Telleschi et al. 2005 for X-ray lines) one finds the following decay laws for chromospheric, transition region UV, transition region FUV, EUV (including the softest X-rays at 20 100 Å), and soft X-rays (1 20 Å), where we also use the rotation law given by Ayres (1997), namely Ω P 1 t 0.6±0.1 for a solar analog (P is the stellar

3 Relative flux (Sun=1) 1000 100 10 1 20 Å 20 100 Å 100 360 Å 360 920 Å 920 1200 Å 1 0.1 0.1 0.2 0.3 0.40.5 1 2 3 4 5 6 7 8 Age (Gyr) Figure 1. Power-law decays in time for various spectral ranges, normalized to the present-day solar flux. Note that the hardest emission decays fastest (from Ribas et al. 2005, reproduced by permission of AAS). rotation period, Ω the angular rotation rate, and t the stellar age): L ch P 1.25±0.15 t 0.75±0.1 (1) L UV P 1.60±0.15 t 1.0 ±0.1 (2) L FUV P 1.40 t 0.85 (3) L EUV P 2.0 t 1.2 (4) L X P 3.2 t 1.9 (5) The heating of planetary upper atmospheres (thermospheres) and the generation of ionospheres is primarily due to the solar EUV to X-ray (XUV) radiation. As the mean free path in the outer atmosphere (in the exospheric layer) becomes sufficiently small, the lighter, heated species may escape into space if their thermal velocity exceeds the gravitational escape velocity (e.g., Kulikov et al. 2007). For sufficiently high gas temperatures, the upper atmosphere may move off hydrodynamically (see, e.g., Watson et al. 1981). Atmospheric evaporation induced by short-wavelength heating may have been responsible for the loss of a large reservoir initially present on Venus. A water oceans would evaporate as a consequence of the strong insolation, inducing a runaway greenhouse in which water vapor would become the major constituent of the atmosphere (Ingersoll 1969). Water vapor was then photodissociated in the upper atmosphere by enhanced solar EUV irradiation (Kasting & Pollack 1983), followed by the escape of hydrogen into space (Watson et al. 1981; Kasting & Pollack 1983). Detailed model calculations suggest that the amount of water in the present terrestrial ocean could have escaped from Venus in only 50 Myr if the H 2 O mass fraction (mass mixing ratio) was at least 0.46, and the XUV flux was 70 100 times the present values (Kulikov et al. 2007). Evaporative losses may also have been relevant for young Earth and Mars although the

4 protecting magnetosphere of Earth (and some magnetism also on Mars) may partly account for a different scenario compared to Venus, and hydrodynamic blow-off may not have occurred (Lammer et al. 2003). Atmospheric evaporation driven by short-wavelength irradiation should then be very important for exoplanets in close orbit around their host stars. Penz et al. (2007) considered Roche-lobe effects to the hydrodynamic loss, and used realistic heating efficiencies and IR cooling terms for HD 209458b, a Jupiter-sized planet in orbit around an old, inactive solar analog. The atmospheric temperatures in their model reach 8000 K at a distance of 1.5R planet which leads to blow-off thanks to gravitational effects by the Roche lobe. These authors obtained atmospheric loss rates of 3.5 10 10 g s 1. Assuming 100 times higher XUV flux at an age of 0.1 Gyr, the loss rate may exceed 10 12 g s 1, leading to an integrated hydrogen mass loss of 1.8 4.4% of the present mass over the entire MS life time of the host star. This loss rate is not leading to substantial alterations in the planet but may still play a significant role in the evolution of its atmosphere. Also, higher evaporation fractions are possible for planets orbiting closer to the host star. In this context, the very slow evolution of magnetic activity in M dwarfs is potentially important for planets in their habitable zones. The high L X (relative to L bol defining the habitable zone) affects planetary atmospheres (and potentially the presence of water) for a much longer time than for a more massive star (e.g., Scalo et al. 2007, see below). 3. High-Energy Environments Apart from (quasi-)steady short-wavelength radiation from magnetically active (chromospheric and coronal) regions on the star, explosive energy release in flares may alter the radiative environment of planets on time scales of minutes to hours, by orders of magnitude in flux. Flares are most likely a consequence of magnetic reconnection in the stellar corona. This process leads to both heating and accelerated electrons which travel toward lower layers where they impact in the chromosphere, thus heating this material and evaporating it into the corona where high levels of X-rays are emitted. Some flares, in particular the most energetic examples, are accompanied by Coronal Mass Ejections (CMEs) that join the stellar wind at high velocities. Extreme flare events may be effective in modifying planetary atmospheres. Schaefer et al. (2000) identified, regardless of main-sequence age, superflares for solar-like stars with total radiated energies of order 10 35 10 38 erg (in X-rays or optical bands). Their recurrence time must be long (decades to centuries), but the irradiation by some of these flares may exceed the total irradiation input by the whole star for maybe an hour. The result could be temporary excess heating, break-up of the ionosphere, and build-up of nitrous oxides at high altitudes. Nitrous oxides in turn destroy ozone; Schaefer et al. (2000) estimate that an event with 10 36 erg of ionizing energy results in 80% of ozone loss for more than a year, thus increasing the ultraviolet irradiation of the planetary surface from normal stellar emission. Again, the prolonged activity of M dwarfs makes this an important issue for planetary habitability.

However, the class of large flares may, on average and on long time scales, not be energetically important for the stellar XUV output. Many flares are too small to be recognized individually in disk-integrated light curves, but these are the most frequent events. Both in the case of the Sun and active stars, X-ray and EUV studies show that flares are distributed in energy according to a power law, dn de = ke α (6) where dn is the number of flares per unit time with a total energy in the interval [E,E + de], and k is a constant. If α 2, then the energy integration (for a given time interval, E max E min E[dN/dE]dE) diverges for E min 0, that is, if the power law is extrapolated to small flare energies, a lower cut-off is required for the power-law distribution. On the other hand, any energy release power is possible depending on the value of E min. Specifically, for young, magnetically active G M dwarf stars including premain sequence objects, monitoring studies in the EUV and X-ray bands found α mostly in the range of 2 2.5 for G M dwarfs (Audard et al. 2000; Kashyap et al. 2002; Güdel et al. 2003; Stelzer et al. 2007), supporting the view that moderate flares are the dominant heating source for these active coronae, and therefore the dominant contributors to short-wavelength radiation. Supporting evidence comes from a linear correlation between the rate of EUV flares (above a given threshold) and the time-averaged X-ray luminosity for late-type stars (Audard et al. 2000), and probably also for T Tauri stars (Stelzer et al. 2007). Why is it important that short-wavelength radiation is due to flares rather than a more steady source? One important aspect for the evolution of planetary atmospheres is the short-term fluctuation of the input energy. Although details are not well studied, fluctuating XUV radiation may drive planetary atmospheres significantly out of equilibrium as a consequence of variable photochemical reactions (e.g., Scalo et al. 2007). Further, although the soft X-ray output of a frequently flaring corona may be nearly equivalent to a constant radiator, the spectrum of the former is different as it should be accompanied by a tail of non-thermal X-rays above 10 kev (Güdel 2009), as a result of the bombardment of the chromospheric layers with accelerated electrons. Flare-produced gamma-rays and X-rays will usually not propagate to the surface of planets with any substantial atmospheres, but a fraction of the hard radiation can be reprocessed and re-emitted as UV light that showers the planetary surface, perhaps at biologically relevant doses (Smith et al. 2004). This secondary radiation may involve damaging effects on living cells, but may also act as an evolutionary driver. 5 4. Stellar winds and non-thermal escape While the solar wind has been well investigated in situ, relatively little is known about such winds from other solar-like stars. None of these winds has so far been detected directly (e.g., from their bremsstrahlung, charge exchange etc), but interesting upper limits to the mass loss rate have been derived from corresponding searches (e.g., Ṁ < 1.7 10 11 M yr 1 for the evolved subgiant Procyon, Drake et al. 1993; Ṁ < 7 10 12 M yr 1 or Ṁ < 3 10 13 M yr 1

6 Figure 2. Interaction between a stellar wind or a coronal mass ejection with a moderate planetary magnetosphere but extended (heated) upper atmosphere. The atmosphere above the compressed magnetospause can be eroded (from Lammer et al. 2007). for the old M dwarf Proxima Centauri, Lim et al. 1996 and Wargelin & Drake 2002; Ṁ < 4 5 10 11 M yr 1 for nearby active solar analogs, Gaidos et al. 2000). The upper limits reported by Gaidos et al. (2000) for solar analogs imply a maximum mass loss during solar history of < 6% of M. Wood et al. (2002) and Wood et al. (2005) discuss a promising indirect approach making use of Lyα absorption in so-called astrospheres. The latter are formed where the outflowing stellar wind collides with the interstellar medium. The heliosphere is permeated by interstellar Hi with T (2 4) 10 4 K (Wood et al. 2002). Much of this gas is piled up between the heliospause and the bow shock, forming the so-called hydrogen wall that can be detected as an absorption signature in the Lyα line. The measured absorption depths in Lyα are compared to hydrodynamic model simulations (Wood et al. 2002, 2005). The amount of astrospheric absorption should scale with the wind ram pressure, P w ṀwV w, where V w is the (unknown) wind velocity (Wood & Linsky 1998). The latter is usually assumed to be the same as the solar wind speed. From nearby stars, the following relations are found for the mass loss rate ( M w ), Ṁ w F 1.34±0.18 X (7) Ṁ w t 2.33±0.55 (8) (Wood et al. 2005). Stellar ionized (and magnetized) winds interact with the upper planetary atmosphere/ionospheres and their magnetospheres (if present). Such interactions further deteriorate the atmospheres as ions are carried away. Important mechanisms include ion pick-up, sputtering, ionospheric outflow, and also dissociative recombination; the ionizing source for the upper planetary atmosphere is the EUV and X-ray emission from the Sun (for a summary see, e.g., Chassefière & Leblanc 2004; Lundin et al. 2007). Such non-thermal processes are

7 Figure 3. Energy spectrum of cosmic rays reaching the top of the atmosphere for present-day Earth (solid), a tidally locked, Earth-like planet at 0.2 AU around a 0.5M star (dashed), compared to the incident spectrum outside the magnetosphere (dash-dotted; from Grießmeier et al. 2005). thought to have fundamentally altered the atmospheres of Venus, Earth, and Mars, leading to strong loss of oxygen (Lammer et al. 2006) and therefore water (after photodissociation of water vapor in the upper atmosphere) from young Venus and Mars. Clearly, a strong magnetosphere as in the case of the Earth shields the lower atmospheric layers; such shielding may have been important in retaining much of the terrestrial water. The amount of shielding obviously depends on the strength of the magnetosphere but also on the strength of the wind and its ability to compress the magnetosphere (Fig. 2). Hot Jupiters or planets in the close-in habitable zones around lowmass M dwarfs should be in tidally locked rotation. The slow rotation produces comparatively small magnetic moments (through dynamo action), and the ram pressure of the stellar wind can compress the magnetosphere sufficiently to expose the upper atmosphere; this effect is enhanced if the upper atmosphere is heated and expanded by strong XUV irradiation. Magnetic protection of the atmosphere is thus strongly reduced (Grießmeier et al. 2005; Lammer et al. 2007), leading to increased erosion, while the cosmic ray flux at the top of the atmosphere will be enhanced given the weaker and smaller magnetosphere (Fig. 3, Grießmeier et al. 2005). We already alluded to coronal mass ejections in the previous section. If such CMEs are ejected at a rate of several per day, they will essentially act like an enhanced solar wind, which will further compress planetary magnetospheres and erode the upper atmospheres (Khodachenko et al. 2007). The combination of enhanced exospheric heating by EUV irradiation with consequent exospheric expansion and a CME wind could induce atmospheric losses of up to tens or hundreds of bars for close-in planets. This may again be particularly important for planets in the habitable zones around M dwarfs (Lammer et al. 2007).

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