Pathways to Earth-like atmospheres

Transcription

1 Orig Life Evol Biosph manuscript No. (will be inserted by the editor) Pathways to Earth-like atmospheres H. Lammer K. G. Kislyakova P. Odert L. Leitzinger R. Schwarz E. Pilat-Lohinger Yu. N. Kulikov M. L. Khodachenko M. Güdel A. Hanslmeier Received: date / Accepted: date Abstract We discuss the evolution of the atmosphere of early Earth and terrestrial exoplanets which may sustain liquid water oceans and continents where life may originate. It is shown that the early atmosphere evolution of Earth-like planets depends mainly on the planet s initial volatile contents of the formation process and is strongly related to a complex interplay of the systems impact history, the weathering time-scales of CO 2 into carbonates, as well as the young host star s EUV radiation and plasma outflow induced atmospheric escape processes. The efficiency of the latter depends on the stellar spectral type and is much more intense when the young star arrived at the Zero-Age-Main-Sequence (ZAMS). Depending on the formation age of the planet, its mass and size, as well as the life-time on the EUV-saturated early phase of its host star, many terrestrial planets may not get rid off their protoatmospheres and could end as water worlds with CO 2 atmospheres and hydrogen- or oxygen-rich -rich upper atmospheres. If an atmosphere of a terrestrial planet evolves to a N 2 -rich atmosphere too early in its lifetime the atmosphere may be lost. We show that the initial conditions set by the formation of terrestrial planets and the evolution of the host star s EUV and plasma envi- H. Lammer, M. L. Khodachenko Austrian Academy of Sciences, Space Research Institute, Schmiedlstr. 6, A-8042 Graz, Austria helmut.lammer@oeaw.ac.at K. G. Kislyakova N.I. Lobachevsky State University, University of Nizhnij Novgorod, 23 Prospekt Gagarina, Nizhnij Novgorod, Russian Federation, and Institute for Physics/IGAM, University of Graz, Universitätsplatz 5, A-8010 Graz, Austria P. Odert, L. Leitzinger, A. Hanslmeier Institute for Physics/IGAM, University of Graz, Universitätsplatz 5, A-8010 Graz, Austria R. Schwarz, E. Pilat-Lohinger, M. Güdel Institute for Astronomy, University of Vienna, Türkenschanzstr. 17, A-1180, Austria Yu. N. Kulikov Polar Geophysical Institute, Russian Academy of Sciences, Khalturina Str. 15, Murmansk, Russian Federation

2 2 H. Lammer et al. ronment are very relevant factors that planets may evolve to habitable worlds. Finally we present an idea how one can test and study the discussed atmosphere evolution hypotheses by future UV transit observations of terrestrial exoplanets within orbits of dwarf stars. Keywords atmosphere formation young stars early Earth habitability 1 Introduction The current atmospheres and their isotopic ratios have resulted from the action of impact delivery and possible accumulation of nebular gases (H, H 2, He) during protoplanetary growth, catastrophic release of volatiles and atmosphere formation after the magma ocean solidification, EUV (λ nm) and plasma driven thermal and non-thermal loss processes of volatiles, impact erosion, mantle outgassing during later evolutionary epochs, as well as anorganically or organically enhanced carbonate formation due to CO 2 weathering processes, (e.g., Kasting, et al. 1984; Kempe and Degens 1985; Matsui and Abe 1986; Zahnle, et al. 1988; Pepin 1991; Chassefière 1996a; 1996b; Abe 1997; Solomatov 2000; Dauphas 2003; Halliday, 2003; Tian, et al. 2005; Kulikov, et al. 2006; Albarède and Blichert-Toft 2007; Lundin, et al. 2007; Elkins-Tanton 2008; Lammer, et al. 2008, 2009a; 2009b; Tian, et al. 2008; Tian 2009; Hanslmeier 2010; Lichtenegger, et al. 2010; Elkins-Tanton 2011; Lunine, et al. 2011). The aim of this work is to discuss briefly the origin and early evolution of atmospheres of Earth-type planets during the active phase of their host stars and discuss scenarios how planets like the Earth could obtain a N 2 -rich atmosphere. In Sect. 2 a brief overview on the formation of Earth-type planets and the influence of giant planets to their initial H 2 O inventory is given. In Sect. 3 we discuss the formation hypotheses of protoatmospheres, including the growth of early atmospheres during the planet s magma ocean solidification process. In Sect. 4 we discuss briefly the role of giant impacts to atmospheric loss and heating. Atmospheric escape by the stellar radiation and plasma environment of the young central star will be addressed in Sect. 5. Finally in Sect. 6 we discuss how we may test atmospheric evolution hypotheses via future UV observations of transiting Earth-type exoplanets orbiting M-dwarfs. 2 Formation of terrestrial planets After the gravitational collapse of an interstellar cloud a star is formed in the center of an extended rotating nebular disk. After the disk reaches a state of thermal and dynamic equilibrium, dust grains grow in the disk through thermal collisions and sediments towards the equatorial plane of a dust layer where planetesimals can originate (e.g., Kusaka, et al. 1970; Nakagawa, et al. 1981). From that stage

3 Pathways to Earth-like atmospheres 3 on, the planetesimals collide and form larger bodies and planetary embryos. Protoplanets continue to grow through the capture and collisions of large planetesimals and planetary embryos. Initial conditions assumed in terrestrial planet formation reflect the state of the protoplanetary disk at the end of oligarchic growth (e.g., Wetherill 1990; Raymond, et al. 2004; Lunine, et al. 2011). In such simulations, planets with masses twice that of Mars to Super-Earths of 4M Earth are usually formed. The composition of these planets range from completely dry to water worlds, where the initial H 2 O content mainly depends on the configuration of the protoplanetary disk and the location of the snowline (e.g., Raymond, et al. 2004; Leinhardt and Richardson, 2005; Lunine, et al. 2011). Terrestrial planets form during and after giant planet migration (Mandell, et al. 2007). A giant planet accretes rapidly all bodies that cross its orbit. As a consequence, the probability that a short periodic comet would hit the Earth over its dynamical lifetime is very low and of the order 1:10 6 but could be very different in planetary systems without gas giants. Comets, which are planetesimals formed beyond the snowline, carried <10% of the total amount of H 2 O to the Earth (Morbidelli, et al. 2000). This is consistent with the D/H isotope ratio of Earth s water inventory which is more asteroidal than cometary (Genda, et al. 2008). The collision probability of comets with the Earth would be orders of magnitude higher if Jupiter and Saturn would be replaced by lower mass planets, like Uranus or Neptune (Morbidelli, et al. 2000). Therefore, in systems like ours without gas giants, comets could be a dominant source for an additional delivery of H 2 O to terrestrial planets within orbits inside the habitable zone (HZ). 3 Formation of protoatmospheres The discussion of the earliest atmospheres can be mainly focused on two cases. All of them have one thing in common, they will produce very dense hydrogen-rich atmospheres. The first possibility would be a hydrogen- and He-rich protoatmosphere which originated from nebula gas. The second case corresponds to a dense water-dominated H 2 O/CO 2 steam atmosphere which is catastrophically outgassed from the magma ocean during its solidification. One should note that both cases are also strongly linked to the impact history of the system. Because the origin of a planet s protoatmosphere determines the initial conditions and further evolution we will briefly discuss the main aspects of these hypotheses in the following subsections. 3.1 Nebular-based protoatmospheres That the early Earth or similar terrestrial planets may have been surrounded by large nebular atmospheres was considered by several researchers because this hypothesis could explain why large amounts of nebular gases such as 3 He, solar-like Ne, or solar-like Ar can be found in Earth s plume basalts (e.g., Clarke, et al. 1969; Mamyrin, et al., 1969; Craig and Lupton 1976; Sasaki and Nakazawa 1988; Pepin, 2000; Moreira, et al. 2001). In case the proto-earth was surrounded by a dense hydrogen-he-type protonebular gas envelope, the thermal effects on the planet s surface would have been dramatic and the surface temperature T s could

4 4 H. Lammer et al. reach >4000 K (Hayashi, et al. 1979; Matsui and Abe 1986). These atmospheric temperatures would have resulted in high temperatures in the outer mantle so that magma oceans could originate and a significant amount of nobel gases with a solar composition could be incorporated in the planet s interior. On the other hand there is much evidence from Xe isotopes that orders of magnitude more have been lost to space. The opponents of the nebular-based protoatmosphere hypothesis expect that the accumulated hydrogen envelopes escape before the planet finished its accretion and the lifetime of gas around stars in the process of planetary formation is most likely limited to a few Myr (Halliday 2003; Najita, et al. 2007; Lunine, et al. 2011). To attract such an amount of gas before the remains of the nebular gas are accreted into the Sun/star or other planetary bodies, in the case of the Earth the time scales should have been in the order of 1 10 Myr (Halliday 2003). Such timescales are very short and it is expected that one would acquire a large nebular atmosphere with nebular timescales of Myr or at least comparable to the accretion phase of the Earth at Myr which can be dated by tungsten and other lead isotopes (e.g., Allègre, et al. 1995; Podosek, et al. 2003; Qingzhu, et al. 2003; Touboul, et al. 2007). As mentiond before, such time scales for nebulae are not in agreement with observations and most disks disappear on time scales of 3 Myr. Disks with ages of 10Myr are rather exceptional. To overcome the before mentioned problem of a nebular-based protoatmosphere for the explanation of Earth s solar-noble gas content one can suggest that these noble gases where incorporated into the Earth from accretion of earlier objects (Farley and Poreda 1992) such as planetary embryos discussed in Sect. 2. On the other hand these planetesimals and embryos could also be fractionated and strongly degassed before they contributed to Earth s formation. Another interesting idea that solar-noble gases might be explained by drop of them into early formed planetesimals is suggested by Podosek (2003). With the evidence that the young Sun was much more active compared to the modern Sun (e.g., Zahnle and Walker 1982; Güdel, et al. 1997; Ribas, et al. 2005, Güdel 2007), this hypothesis gets a strong support, because one can expect that inner Solar System planetesimals should have incorporated many solar-wind implanted noble gases. However, one can not completely rule out an accumulation of hydrogen-rich nebular-based remnants of gaseous envelopes around early Earth and similar planets. Hayashi et al. (1979) and Mizuno et al. (1980) suggested that in stages earlier than the formation of Jupiter or gas giants, rocky protoplanets grew in the gaseous disk and accumulated hydrogen-rich envelopes which became more and more massive with the growth of the protoplanets. After the nebular gas was blown away the protoplanets and planetary embryos which formed the early Earth or other terrestrial planets may not have completely lost these gaseous envelopes by EUVpowered blow-off and the accreted planet could be surrounded by a hydrogen-rich remnant of nebular gas. 3.2 Outgassed and impact-induced protoatmospheres: atmospheric growth during magma ocean solidification That the early Earth and similar planets outgassed their atmospheres is based on hypotheses which were suggested decades ago by Brown (1949) and Rubey (1951).

5 Pathways to Earth-like atmospheres 5 This theory is backed mainly by the similarity in the chemical composition between the atmosphere and hydrosphere with the composition of volcanic gases. Halliday (2003) reviewed several models for the history of degassing of the early Earth which are based on the hypothesis that a relatively undegassed lower mantle supplied the upper mantle with volatiles which were released into the atmosphere. This hypothesis is supported by the isotopic composition and concentration of 40 Ar which provides an evidence that gases have been added to the atmosphere from the planet s interior over geological timescales (e.g., Porcelli and Wasserburg 1995; Allègre, et al. 1996). From Xe isotope studies related to the Earth mantle and atmosphere it was found that the rocky body and the atmosphere separated from each other during an early stage (e.g., Allègre, et al. 1996; Halliday 2003). However, it is highly possible that a fraction of the Xe in Earth s early atmosphere was added after the degassing process via impacts (e.g., Javoy 1998; Owen and Bar-Nun 1995; Morbidelli, et al. 2000; Dauphas 2003) so that differences between elemental and atomic abundances of mantle noble gases relative to the atmosphere can be explained by heterogenous atmospheric accretion (e.g., Marty 1989; Caffee, et al. 1999; Dauphas 2003). That impactors besides the usual accretion process played an important role for the addition of H 2 O can also be considered (e.g., Anders and Owen, 1977; Zahnle, et al. 1988; Lunine, et al. 2003; Matsui and Abe 1986; Morbidelli, et al. 2000; Albarède and Blicher-Toft, 2007; Lunine, et al. 2011). In Earth s case the buoyant hydrous phases were produced by the magma ocean and icy impactors which originated from beyond the snowline and kept it permanently molten or molten in a series of molten phases for a few tens of Myr (e.g., Albarède and Blicher-Toft 2007). Depending on the initial volatile inventories which can be accreted during impacts of planetesimals and embryos discussed in Sects. 2 and based on meteorite studies a wide range of initial H 2 O contents between weight percent (wt%) in the bulk silicate of terrestrial planets can be expected (e.g., Jarosewich 1990; Liu 2004; Elkins-Tanton 2008; Elkins-Tanton 2011). The pressure gradient in the magma ocean of the accreted planet controls solidification and is related to its mass. Because the majority of the minerals that solidify from a bulk planetary magma composition are denser compared to the coexisting magma, the solidification process proceeds from the bottom upward to the surface with increasing amounts of magma lying above solid silicate minerals (e.g., Solomatov 2000; Albarède and Blicher-Toft 2007; Elkins-Tanton 2008; Elkins-Tanton 2011). The mantle solidification is a fast process and more or less complete in 3 Myr for magma ocean depths 2000 km and 10 5 years for low-volatile magma oceans (e.g., Elkins-Tanton 2008) H 2 O and CO 2 molecules can enter solidifying minerals only in small quantities (Elkins-Tanton 2008). As a result they will degas into early water vapour dominated dense H 2 O/CO 2 -rich steam atmospheres after Earth s accretion (e.g., Liu 2004; Elkins-Tanton 2008; Elkins-Tanton 2011). In case that all the expected surface and upper mantle H 2 O and CO 2 could be converted into a primitive atmosphere, an atmosphere with 560 bar of H 2 O and about 100 bar of CO 2 could build-up (Elkins-Tanton 2008). For higher initial volatile contents, these models can produce early atmospheres with surface pressures up to several 10 4 bar (see Table 1). Under such extreme condition supercritical water can originate directly from progressive solidification of a magma

6 6 H. Lammer et al. Table 1 Surface pressure of outgassed water dominated steam atmospheres in units of bar and in Earth Ocean s (EOs) for a planet with 1M Earth and various initial bulk H 2 O inventories inside a 2000 km deep magma ocean in weight percent (wt%) (Elkins-Tanton 2008; Elkins- Tanton 2011) 0.05 wt % 0.1 wt% 0.5 wt % 1 wt % 3 wt % 1M Earth 250 bar 500 bar 7000 bar bar bar 1 EO 2 EOs 28 EOs 60 EOs 92 EOs ocean if the initial H 2 O inventory is above 1 3 wt%, which may be possible for Super-Earths (Elkins-Tanton 2011). In the supercritical phase liquid H 2 O expands and molecules become steam and rise above the liquid. The vapor density consequently increases, while liquid density decreases. When the pressure reaches 220 bar, the liquid and vapor densities are similar and become supercritical - an environment where H 2 O exist in liquid and gas form together. For an Earth-like planet, surface temperatures near the supercritical point of H 2 O can be reached within several millions to tens of Myr after the solidification of a magma ocean. These timescales agree also with studies by several researchers who investigated the early stages of accretion and impacts and expect that due to thermal blanketing temperatures 1500 K could keep the volatiles in vapour phase within several tens of Myr or even up to 100 Myr (e.g., Hayashi, et al. 1979; Mizuno, et al. 1980; Matsui and Abe 1986; Zahnle, et al. 1988; Abe 1997; Albarède and Blichert-Toft 2007). In case that the entire magma ocean solidification process, overturn, and cooling to clement conditions occurs within a faster time of 5 10 Myr, Elkins- Tanton (2008) points out that such a short cooling interval indicates that serial magma oceans of varying depths were probably unavoidable on early terrestrial planets. At the end of such epochs the water vapour which did not escape to space can condense and form oceans. For Super-Earths with dense H 2 O steam atmospheres up to several 10 4 bar this timescale may last a few hundreds of Myr (Elkins-Tanton 2008). In the case of an Earth-like planet, after the surface temperature of the young planet cools below a critical point of 650 K, corresponding to a pressure of 220 bar, the supercritical state collapse into a liquid surface water ocean (Elkins-Tanton 2011). For Earth Elkins-Tanton (2008) suggests that initial magma ocean water inventories 0.1 wt% were unlikely, because they would have implied initial atmosphere pressures of thousands of bars and times related to cooling of clement surface conditions may have lasted hundreds of Myr. Such long cooling times in combination with high surface pressures do not agree with the evidence for liquid H 2 O on Earth s surface discovered in zircons 4.4 Gyr ago (Valley, et al. 2002). If Earth would have obtained slightly more material from water-rich planetesimals, its CO 2 content would have been much higher and the ocean would have been hundreds of kilometers deep. This would have resulted in a globally covered water world which is surrounded by a dense CO 2 atmosphere. The related greenhouse effect would have produced a hot evaporating upper ocean layer. Thus, from the results of these studies we can conclude that a wide variety of initial atmospheres with different water and CO 2 contents should be possible. For the evolution of an Earth-like atmosphere and habitat this is crucial because many planets even if they orbit within the HZ of their host stars may not get rid

7 Pathways to Earth-like atmospheres 7 off their dense protoatmospheres and may end as water worlds or hydrogen-rich sub-uranus-type bodies. This presumption is supported by observation of the two Super-Earths Kepler 11b and Kepler 11f with 4.3M Earth and 2.3M Earth which have radii of 1.97R Earth and 2.61M Earth and mean densities of 3.1 and 0.7 g cm 3 (e.g. Borucki, et al. 2011; Lissauer, et al. 2011). These densities indicate substantial envelopes of light gases such as H/He or H 2 O/H. 4 Giant impacts and their influence on early atmospheres During planetary formation several tens of Mars-sized planetary embryos can originate after the successive accretion of planetesimals (e.g., Kokubo and Ida 1998; Raymond, et al. 2004). These bodies would have an impact-induced, a solar-type or a mixed atmosphere (Zahnle, et al. 1988) and collide with each other (Chambers and Wetherill 1998). A few Myr after Earth finished its accretion, a giant impact occurred most likely with a Mars-sized body which resulted in the formation of the Moon (e.g., Touboul, et al. 2007). Because co-accretion (Taylor 1992) and capture (Urey 1966) cannot explain the angular momentum of the Earth-Moon system, and because of the difference in density and volatile depletion and similar O isotope compositions it is expected that this giant collision ejected a cloud of fragments, which aggregated into a full or partial ring around Earth and then coalesced into a proto-moon (Hartmann and Davis 1975). Kleine et al. (2002) suggests from small tungsten isotopic effects of the Moon, that it was formed Myr after the origin of the Sun s origin which would indicate an age of the Earth of Myr. In a recent study Touboul et al. (2007) found a new constraint for the core formation of the Moon at around 62 Myr from a W isotope analysis of lunar rocks. This new constraint provides the earliest time that the giant impact could have formed the Moon and corresponds to an accretion age of the early Earth in agreement with Allègre et al. (1995) of 50 Myr after the origin of the Sun. When such a giant impact occurs, a part of the target atmosphere escapes due to global motion excited by a strong shock wave traveling to the planetary surface and its interior. The fraction of the lost atmosphere to the total mass of the atmosphere is found to be controlled mainly by the velocity at the planetary surface. By investigating the escape fraction of an atmosphere by a Moon-forming giant impact early studies estimated that Earth s primordial atmosphere was completely lost (Ahrens 1993; Chen and Ahrens 1997). Modern impact studies indicate that only 30% escape and most of the atmosphere survives (Genda and Abe 2003). We note that these studies assumed a hydrostatically equilibrated polytropic atmosphere. This restriction may lead to inaccurate results, because as discussed in the following section due to the active young Sun/star early atmospheres were/are most likely non-hydrostatic. As discussed in Sect. 3 and according to several accretion and magma ocean studies, after the formation of a dense water vapour dominated H 2 O/CO 2 steam atmosphere it can take a few tens of Myr that the surface temperature cools down to the temperature range that the critical pressure level is reached and H 2 O can rain out and build a liquid ocean on a planetary surface (e.g., Mizuno, et al. 1980; Liu 2004; Albarède and Blichert-Toft 2007; Elkins-Tanton 2008; Elkins-Tanton

8 8 H. Lammer et al. 2011). However, the before mentioned studies neglect the possible cooling in a convective dense atmosphere (Zahnle, et al. 1988). On the other hand frequent impacts, which are expected during the first tens of Myr after the origin of the Solar System and thermal blanketing may have kept the environment of Venus and Earth hot enough during several tens of Myr (e.g., Hayashi, et al. 1979; Mizuno, et al. 1980; Matsui and Abe 1986; Abe 1997; Albarède and Blichert-Toft 2007). Impacts of large bodies will vaporize and dissociate a huge fraction of H 2 O and release CO 2 from surface rocks into the atmosphere. These effects lead to an efficient temporary greenhouse effect so that the atmosphere can be heated to temperatures which are hot enough that rocks on the planets surface can be melted. After the impact phase was over and the impact energy rate decreased to negligible values, the equilibrium thermal structure for a planet such as the Earth can reach the critical point that the water can condense to a liquid. In the cases such as that of Venus, due to the planet s closer orbital distance to the Sun the temperature is higher so that the vapour remains in its gaseous form. 5 The role of the radiation environment of young stars on atmospheric evolution The evolution of planetary atmospheres can only be studied together with the history of the host stars activity (e.g., Lammer, et al. 2008). The Xe isotope data provide evidence that Earth lost a huge amount of its early atmosphere within the first 100 Myr after the planet s origin (e.g., Ozima and Podosek 1999; Porcelli, et al. 2001; Halliday 2003). Due to the high EUV flux of the young Sun/star and frequent impacts the H 2 O molecules in the thermosphere are dissociated and H 2 and H atoms will dominate the upper atmosphere as long as they are not lost in a combination between thermal and nonthermal escape processes and giant impacts. Depending on atmospheric composition and planetary mass, when the EUV flux in the wavelength range λ nm) exceeds a critical value, the outward flow of the bulk upper thermosphere cools due to adiabatic expansion (Tian, et al. 2005; 2008). As it was shown by studies of Watson et al. (1981), Kasting and Pollack (1983) and Tian et al. (2005), for a EUV flux which is only a few times larger than that of the present Sun, the exobase of a hydrogen-rich upper atmosphere can expand several planetary radii above the planet s surface. Fig. 1a shows the critical temperature T c (e.g., Öpik 1963; Chamberlain 1963) T c = 2m HGM pl, (1) 3kr where m H is the mass of a hydrogen atom, k the Boltzmann constant, M pl the planetary mass and r planetocentric distance, for atomic hydrogen blow-off as a function of exobase distance. The asterisk symbols mark exobase levels for a hydrogen-rich thermosphere on Earth which is exposed to a present, 2.5 and 5 times higher solar EUV flux. According to Tian et al. (2005) for a 5 times higher EUV flux than that of the modern Sun the exobase level which separates the collision dominated thermosphere form the collision-less exosphere moved up to 7R Earth. Therefore, it can be expected that for EUV fluxes times ( 4.3 Gyr ago) compared to today s Sun, blow-off of hydrogen can certainly occur. In case of the absence of a high amount of IR-cooling molecules in the thermosphere, at such high EUV fluxes even atomic oxygen reached most likely blow-off conditions.

9 Pathways to Earth-like atmospheres Observation based EUV-enhancement factors of young stars a b 10 3 T c [K] 1000 S EUV EUV z exo [km] t [Myr] Fig. 1 a: Critical temperature for hydrogen blow-off on Earth as a function of exobase distance. The symbols indicate exobase levels for a hydrogen-rich thermosphere exposed to the present time solar EUV flux and enhanced fluxes which are 2.5 and 5 times larger (Tian, et al. 2005). b: Solid line shows the EUV-enhancement factor S EUV as obtained form the detailed analysis of solar analogues with different ages (Güdel, et al. 1997; Ribas, et al. 2005). The dotted line shows S EUV as obtained by Zahnle and Walker (1982) from scalings and extrapolations of solar data and a T-Tauri star. In previous studies on EUV-powered hydrodynamic escape of hydrogen from growing protoplanets and early H 2 O steam atmospheres (e.g., Zahnle, et al. 1988) EUV flux enhancement approximations for the young Sun of S EUV = 10 4 /t Myr for t Myr 1 Myr were applied during time periods of Myr after the Sun s origin. Their S EUV was not based on observational data of solar analogue stars with different ages but scaled from solar data and observations of a T-Tauri star (Zahnle and Walker, 1982). In the meantime observations of solar analogue stars with different ages revealed that the EUV flux is saturated at an enhancement value of 100 during the first 100 Myr (Güdel, et al. 1997) and decreases then according to S EUV = (t Gyr /t Sun[Gyr] ) 1.23 (Ribas, et al. 2005; Güdel 2007). Fig. 1b shows a comparison of the EUV-enhancement factors S EUV which were applied in the pioneering works of EUV-powered hydrodynamic escape estimations by Zahnle, et al. (1988) and that inferred from the detailed analysis of solar proxies (Ribas, et al. 2005) by taking into account an EUV-saturation phase during the first 100 Myr (Güdel, et al. 1997). One can see that the S EUV applied by Zahnle, et al. (1988) overestimated the EUV flux of the young Sun during the first 100 Myr up to two orders of magnitude. For time periods between present and 4.4 Gyr ago both S EUV yield more or less similar results. 5.2 Hydrodynamic escape of H atoms from primitive atmospheres Due to the high EUV flux of the young Sun or other stars the H 2 O molecules which populate the atmospheres of early terrestrial planets are dissociated and atomic

10 10 H. Lammer et al. and molecular hydrogen populate the upper atmosphere as long as they are lost to space. Photons in the EUV-range can be absorbed by the light hydrogen atoms and molecules and are uniquely suited to induce thermal escape. By applying the energy-limited 1 escape formula (e.g., Hunten 1993) dm dt = 3ηS EUVF EUV 4Gρ pl, (2) with heating efficiency η, the gravitational constant G, the mean planetary density ρ pl and the present time EUV flux F EUV in the wavelength range nm in 1 AU as shown in Fig. 2a Zahnle et al. (1988) obtained an equivalent loss of hydrogen from early Earth during the protoplanetary stage up to accretion Myr of 25 Earth oceans (EO). From atmospheric escape studies of hydrogen-rich Hot Jupiters it is known that the thermal mass loss rate obtained from the energylimited formula yields overestimated loss rates, but can be well approximated by a modification which includes a mass loss enhancement factor due to a Roche lobe effect and a realistic η 15% for the stellar EUV radiation (Erkaev, et al. 2007; Penz, et al. 2008; Lammer, et al. 2009; Leitzinger, et al. 2011). Because the terrestrial planets considered here are located at larger orbital distances compared to Hot Jupiters we can neglect the Roche lobe effect but have to introduce a realistic η. Lower values for η are also in agreement with previous studies related to the loss of Venus initial water inventory (Watson, et al. 1981; Kasting and Pollack 1983; Chassefière 1996a) and a recent study by Murray-Clay et al. (2009) who modeled the EUV driven mass loss from Hot Jupiters during their host stars early active phase and found that η changes with time. Murry-Clay, et al. (2009) discovered that for high EUV fluxes as expected during the first 100 Myr after the origin of the host star, Lyman-α cooling of the thermosphere is relevant and η should be <20 %. In this process Lyman-α radiation is emitted by atomic a b EO H EO H t [Myr] t [Myr] Fig. 2 Amount of atomic hydrogen in corresponding amounts of Earth Oceans (EOs) that can be lost shortly after the origin of the solar system (a) and after the expected time period of 50 Myr when Earth finished its accretion, based on the two different EUV-enhancement factors (dotted line: Zahnle, et al. 1988; solid line: Ribas, et al. 2005). 1 Energy limited means that the ratio of the net heating rate to the rate of stellar energy absorption η = 100%.

11 Pathways to Earth-like atmospheres 11 hydrogen atoms which are collisionally excited by electrons. Because we are also interested in the escape efficiency of the EUV radiation during the early periods of the planets we consider in agreement with the before mentioned studies an η 15% but an S EUV according to Ribas et al. (2005) and obtain a much lower upper limit for hydrogen escape, equivalent of <7 EOs. The overestimation results from the much higher EUV flux related to the S EUV applied by Zahnle et al. (1988) and the energy-limited approach with a heating efficiency η=100 %. Moreover, the availability of IR-cooling molecules such as H + 3 and CO 2, as well as the dragging of heavy species such as remaining oxygen atoms within the hydrogen outflow would additionally reduce the escape rate (e.g., Zahnle, et al. 1988; Kulikov, et al. 2006). Fig. 2b shows the maximum amounts of H atoms for a heating efficiency η=15% in units of bar after finished accretion of the Earth at 50 Myr (Allègre, et al. 1995; Touboul, et al. 2007) during 450 Myr for both EUV-enhancement factors shown in Fig. 1. The solid line corresponds to the hydrogen content of 6 EOs and the dotted line to 12 EOs. From our escape estimations one can see that the S EUV used in Zahnle et al. (1988) yield an overestimated hydrogen loss of 6 EOs. However, our loss estimations are very conservative because the protoplanet will also lose atmosphere before its accretion is finished. It is expected that about 80 % of Earth s mass was reached 20 Myr and 90 % at 30 Myr after the origin of the Sun (Zahnle, et al. 1988). If a terrestrial planet would produce a water vapour dominated steam atmosphere after its accretion of 1600 bar its hydrogen fraction cannot escape completely during the EUV active phase of a solar like G star and a hydrogen envelope may remain. Furthermore, if EUV-powered hydrogen blow-off is efficient enough that the hydrogen content of several EOs can escape from a planet a huge amount of oxygen should be accumulated. 5.3 Hydrodynamical escape of dragged O atoms The problem that large amounts of oxygen could be accumulated during a strong escape of hydrogen from H 2 O-rich primitive atmospheres was first addressed by Kasting (1995). Stimulated by this argument Chassefière (1996b) investigated the hydrodynamic loss of oxygen from primitive atmospheres of Venus and Mars in detail. However, the study of Chassefière (1996b) is based on terrestrial planet formation models in which the time of the final accretion for the Earth and Venus occurred 100 Myr after the formation of the Sun (Wetherill 1986). At that time the EUV-saturation phase of the young Sun ended and S EUV decreased according to the S EUV power laws discussed before and shown in Fig. 1b. As discussed in Sect. 4 from W isotope studies in lunar metals it is now expected that the Earth was most likely accreted 50 Myr after the Sun s origin and obtained a mass of 80 90% of its present even 30 Myr earlier (e.g., Quingzhu, et al. 2002; Touboul, et al. 2007). These data indicate that Earth most likely finished its formation several tens of Myr earlier than the formation time which was assumed in previous atmospheric escape studies. Moreover, in the study by Chassefière (1996b) the cooling phase of the magma ocean was expected to be >100 Myr, while more recent studies indicate that the solidification of magma oceans with a depth of 2000 km are fast processes and mantle solidification of 98% can be completed in 5 Myr (e.g. Elkins-Tanton

12 12 H. Lammer et al. 2008; Elkins-Tanton 2011). The assumption of the late finished planets resulted in ages where the EUV flux of the young Sun was much lower compared to the EUVsaturation phase. While Kasting and Pollack (1983) studied the hydrodynamic escape of hydrogen and water from early Venus with EUV flux values which can be compared to that of the present Sun, Chassefière (1996b) applied as its highest value an S EUV which was 25 times higher than that of the present Sun. As it is shown from the detailed analysis of solar analogue stars (Güdel, et al. 1997; Ribas, et al. 2005), the evidence of faster formation (e.g., Quingzhu, et al. 2002; Touboul, et al. 2007) and short nebular evaporation times (e.g., Najita, et al. 2007; Lunine, et al. 2011) the accreted or nearly accreted terrestrial planets in the Solar System were most likely exposed during a time period of Myr to an EUV flux which was 100 times higher compared to that of today s Sun. Because of the different initial parameter values which were assumed in the previous studies of Kasting and Pollack (1983) and Chassefière (1996b) and the complex links between the various astrophysical, geophysical and aeronomical factors it is difficult to comment on the accuracy of these early results. It is certainly beyond the scope of this study that we reinvestigate the pioneering works of Kasting and Pollack (1983), Kasting and Zahnle (1986), Zahnle et al. (1988) and Chassefière (1996), but we carry out a rough investigation on the expected consequences of early outgassed water dominated H 2 O/CO 2 steam atmospheres at early Earth when the planet finished the accretion during the EUV-saturated phase of the young Sun. In an atmosphere which is dominated by water vapour, the H 2 O molecules will be dissociated by the EUV and X-rays as well as frequently occurring impacts in the lower thermosphere (Chassefière, 1996b). Therefore one can expect that O atoms will be the major form of oxygen which can be dragged by the outward flowing hydrogen flux F H in the formula given by Hunten et al. (1987) ( F O = X O F H X H ) m H + ktf H bgx H m O ( ) m H + ktf H bgx H m H = X O X H F H f, (3) with the escaping flux F O of atomic oxygen, the mole mixing ratio X H and X O, their masses m H, a molecular diffusion parameter b for O atoms in hydrogen given in Zahnle and Kasting (1986), gravity acceleration g, Boltzmann constant k, and an average upper atmosphere temperature T which can be assumed for hydrogen under such conditions in the order of 500 K (Zahnle and Kasting, 1986; Chassefiére, 1996b). There are basically two conditions that the formula in Eq. (3) is valid (Zahnle and Kasting, 1986; Chassefière 1996b). The first condition requires that the mass of the dragged particle is compared to the main gas. The second condition is linked to the isothermal assumption of Hunten et al. (1987) that the fractionation factor f is > m H /m O, which is in our case also verified (Chassefière 1996b). In our rough estimations we assume that the dragging of the O atoms stops, when the number density becomes equal compared with the hydrogen. Fig. 3a shows the calculated loss of an initial 500 bar ( 2 EOs) dense H 2 O steam atmosphere by using a heating efficiency η=15% and S EUV of 100 during the saturation phase of the young Sun and under the assumption that the surface and atmosphere reached the critical temperature of 650 K where the remaining H 2 O-vapour of 1EO condensed and collapsed into the liquid surface water ocean. One should note that this is a rough assumption and additional water could have

13 Pathways to Earth-like atmospheres a 2.0 b 1.5 O 1.5 EO H,O,H2O 1.0 H, H 2 O EO H,O 1.0 H 0.5 O residual 0.5 O t [Myr] t [Myr] Fig. 3 a: Expected loss scenario of an outgassed 500 bar ( 2 EOs) H 2 O-vapour steam atmosphere from early Earth which finished accretion at 50 Myr. The solid line related to H atoms shows the escape of the hydrogen content from the initially outgassed H 2 O amount until the equivalent amount of 1EO is reached. The dashed line corresponds to the dragged O atoms which accumulate to a residual fraction of a corresponding amount of 0.5EO when the H blow-off stops 25 Myr after the atmosphere outgassing due to the formation of Earth s ocean. The remaining O atoms escape during the following 150 Myr of the high activity phase of the young Sun. b: Shows a similar scenario but at Venus orbit at 0.7 AU where thermal blanketing and the higher and increasing luminosity of the young Sun prevents formations of a liquid ocean due to higher surface temperatures (Matsui and Abe 1986). In such an environment an upper value of 2 EOs of H 2 O vapour can be lost during the EUV-saturation phase of the young Sun/star. been delivered by later impacts too, but the majority of early Earth s initial water inventory could have been produced by condensation of a fraction of the outgassed H 2 O dominated steam atmosphere. Under this assumptions up to the ocean formation the hydrogen fraction would be lost within 25 Myr after the atmospheres origin but an equivalent amount of 0.5 EO accumulated oxygen would remain. If one assumes that these remaining oxygen atoms populate the upper atmosphere they may also experience blow-off conditions and strong nonthermal escape during the following 150 Myr where the solar EUV flux was times higher than today. One should also note that a fraction of 30% of atmosphere could have been also lost during the Moon-forming impact (e.g., Gender and Abe 2003). Moreover, from our estimations in agreement with Elkins-Tanton (2008) and Elkins-Tanton (2011) we can draw two interesting conclusions. First, the early Earth should not have produced a H 2 O dominated steam atmosphere which was much denser than 500 bar because higher amounts could accumulate most likely large oxygen remnants. Second, if the Earth s surface temperature reached the right temperature and pressure conditions for H 2 O condensation ( 650 K, 250 bar) and its collapse into Earth s ocean around Myr after the atmosphere formation, accumulated oxygen with an amount of 0.5 EOs could have been lost during the next hundreds of Myr. If the early Earth would have accumulated much higher amounts of oxygen or produced its liquid ocean later, due to the decreased EUV flux after the saturation phase it may have had a problem for losing it. Fig. 3b shows the water loss of an amount of 500 bar from an Earth-like planet around a Sun-type star in Venus orbit. One can see that in such a scenario the Earth-like planet would have lost the hydrogen amount of 2 EOs H 2 O during

14 14 H. Lammer et al O EO H,O 3.0 H t [Myr] Fig. 4 Escape and accumulation of H and O atoms in EOs from a 1000 bar ( 4 EO) water vapour dominated outgassed steam atmosphere at a 50 Myr old Earth-like planet during the EUV-saturation phase of a solar-like G-type star. 25 Myr with a remaining oxygen content of 0.4 EOs. By assuming that the remaining O atoms experienced also blow-off during the high EUV activity phase of the young Sun they would be lost 30 Myr later. In case that high levels of CO 2 molecules cooled the lower thermosphere to a level that O blow-off was not possible, then a remnant of several tens of bar could be lost by nonthermal escape processes (Kulikov, et al. 2006). If the outgassed H 2 O steam atmosphere is much denser than 500 bar the planet may have also a problem to get rid of dense hydrogen and oxygen envelopes. Fig. 4 shows the escape of hydrogen and oxygen from an initially outgassed 1000 bar dense H 2 O-vapour steam atmosphere during the EUV-saturation phase of the young Sun or a similar G-type star. In such a case an Earth-mass planet keep more than 2.3 EO amounts of hydrogen and more than 3.2 EO amounts of oxygen. That such a huge amount of oxygen can be lost via solar wind related ion erosion and other nonthermal atmospheric escape processes is very unlikely. Our rough estimates agree with the suggestions of Kasting (1995) and Chassefière (1996) that there may be planets depending on their size, mass, orbital distance as well as their host star s EUV flux evolution, which accumulate huge amount of abiotic oxygen. Super-Earths which may outgas water vapour dominated steam atmospheres with several 1000 to several 10 4 bar (Elkins-Tanton 2011) will be good candidates for such atmospheres. 5.4 EUV-response and stability of Earth-like N 2 -rich atmospheres An other problem related to the evolution of early Earth s atmosphere is that its nitrogen inventory should have been protected against atmospheric escape during the first hundreds of Myr after the Earth s formation. Similar as for hydrogenrich thermosheres a N 2 -dominated thermosphere of the early Earth would also experience a fast transition to a hydrodynamic expansion when it was exposed

15 Pathways to Earth-like atmospheres 15 to sufficiently high EUV radiation. If the EUV flux reaches 7 times that of the present Sun, the exobase temperature exceeds K (Tian, et al. 2008; Lammer, et al. 2008). In such a case the exobase could move above the present average subsolar magnetopause stand-off distance of 10R Earth so that the neutral constituents beyond the magnetopause can be ionized and picked up by the solar wind, capable to erode a 1 bar N 2 atmosphere during 10 Myr (see Table 2) (Lundin, et al. 2007; Lichtenegger, et al. 2010). A higher amount of CO 2 between Gyr ago may also confine the upper atmosphere within the shielding magnetosphere and thus could have protected Earth s N 2 inventory from escape. However, model simulations show that an Earth-like planet with even a 95% CO 2 atmosphere should experience problems related to atmospheric escape during the first Myr after the Sun arrived at the ZAMS, when the EUV flux reached values 30 time that of the present Sun (Tian 2009; Lichtenegger, et al. 2010). Fig. 5 Atmosphere evolution scenarios of Earth-type planets after the outgassing of volatiles and fast growth of dense water vapour dominated H 2 O/CO 2 steam atmospheres (e.g., Elkins- Tanton 2008; 2011). A complex interplay between the young star s EUV activity, atmospheric escape processes, large impacts and weathering time-scales of CO 2 into carbonates during the first 500 Myr resulted in the formation of Earth s 0.8 bar N 2 -rich atmosphere Gyr ago (Kempe and Degens 1985; Kasting, et al. 1984). If the outgassed amount of water vapour is too large the oxygen fraction of the H 2 O may only lost partially and an abiotic oxygen-rich atmosphere can accumulate (Chassefière 1996). Earth-like N 2 -dominated atmospheres on early terrestrial are not stable against EUV-induced atmospheric expansion and nonthermal escape as long as the EUV flux is 7 times than that of the present solar value (Lichtenegger, et al. 2010).

16 16 H. Lammer et al. Table 2 Escape of a N 2 -rich Earth-like atmosphere from a planet with the size and mass of the Earth in units of bar as a function of solar/stellar EUV flux and wind exposure time t, for a weak wind with corresponding values as at present Earth t w in 1 AU and a wind t st which is assumed to be 30 times denser (Lichtenegger, et al. 2010) EUV t Earth [Myr] t w =10Myr t w =100Myr t st =10Myr t st =100Myr bar 10 bar 5 bar 50 bar bar 4 bar 3 bar 30 bar bar 2 bar 1 bar 10 bar Moreover, depending on the alkalinity of the early ocean water, CO 2 may be weathered effectively as CaCO 3 so that a substantial amount of CO 2 from the primitive atmosphere could have been removed within a short geological time period so that nitrogen remained (Kempe and Degens, 1985). In such a scenario the nitrogen inventory should have escaped to space effectively (Lichtenegger, et al. 2010). Because Earth s atmosphere is not enriched in 15 N compared to the solar 15 N/ 14 N ratio one can expect that its initial N 2 inventory was kept until today. Based on observed similarities of isotopic fractionation in Earth s and Titan s atmospheres some researchers argue that the late heavy bombardment 3.8 Gyr ago may have enriched Earth s volatile content including the N 2 inventory (Trigo- Rodriguez and Martin-Torres 2011). To overcome this problem a remnant of a dense and extended hydrogen or accumulated oxygen envelope from the initially outgassed H 2 O inventory, which could survive or form after the Moon-forming impact could have also acted as a shield against the loss of heavier atoms and molecules such as the Earth s initial N 2 inventory. If Early Earth s N 2 inventory was indeed protected by such envelopes, such a remnant gas should have been lost until today. Fig. 5 summarizes in an illustration the expected atmosphere evolution scenarios for early Earth, Venus and other terrestrial planets. 6 Testing atmospheric evolution scenarios by UV-transit follow-up observations The detection of EUV heated extended non-hydrostatic upper atmospheres around terrestrial exoplanets would constrain the uncertainties related to the discussions addressed in the previous sections. Since the discoveries of several Super-Earths the search for terrestrial exoplanets within HZs of low mass M stars become feasible. As discussed in the previous sections one can expect that upper atmospheres will expand to several planetary radii if no effective IR-cooling molecules are available in the thermosphere when they are exposed to several times higher EUV fluxes compared to that of the present Sun. As illustrated in Fig. 6 as a consequence the upper atmosphere expands beyond a possible magnetopause obstacle and energetic neutral atoms (ENAs) will be produced via charge exchange collisions between the charged particle flow of the planet s host star. In case we can observe extended upper atmospheres and related ENA clouds around Earth-type exoplanets, we would obtain information about the stellar wind

17 Pathways to Earth-like atmospheres 17 plasma properties, the upper atmosphere structure, non-thermal ion escape, and the planet s magnetic or non-magnetic obstacle. For the simulation of the ENA production, one can apply plasma flow models or particle flow models in combination with upper atmosphere and exosphere models (Holmström, et al. 2008; Ekenbäck, et al. 2010; Lammer, et al. 2011a). For illustrating an expected stellar wind interaction between an Earth-size exoplanet with a hydrogen-rich upper atmosphere which is exposed with a 5 times higher EUV flux compared to the modern Sun we apply as in Ekenbäck et al. (2010) and Lammer et al. (2011a;b) a direct simulation Monte Carlo code which is coupled with a stellar wind plasma flow model and obtain the preliminary results shown in Fig. 7. In that case the exobase level expands to about 7R Earth and we exposed the huge planetary hydrogen corona with a stellar wind density of 500 cm 3, and a velocity of 300 km s 1. Such plasma parameters can be expected within the HZs of moderate active M-type stars. Terrestrial planets around M-stars would be promising candidates for such observations because most of them are much longer active in the EUV range compared to Sun-type G-stars and their HZs are located in orbital distances which are <2 AU. One can expect dense plasma environments comparable with the early solar wind (Lammer, et al. 2011b). After the discovery of a transit of a terrestrial exoplanet, follow-up observations within the Lyman-α range for the detection of emission absorption caused by extended hydrogen and ENA clouds with the upcoming World Space Observatory- UV (WSO-UV) could be performed. The WSO-UV is a Russian-led international space telescope devoted to high resolution UV spectroscopy and imaging which is included in the Federal Space Program of the Russian Federation with a foreseen launch date around (Shustov, et al. 2009). The confirmation of an observation of a hydrogen coronae and/or ENA clouds in Lyman-α during the transit of a rocky exoplanet would indicate that its upper atmosphere is in a nonhydrostatic stage. Such observations which can be expected in the not so distant future will have an important impact in our understanding on the evolution of terrestrial planet atmospheres, including that of early Venus, Earth and Mars. Fig. 6 Left, a hydrostatic atmosphere like that of present Earth which is protected by its magnetosphere. Right, illustration of a EUV-heated upper atmosphere which extends above the magnetopause, so that energetic neutral atoms (ENAs) will be produced via charge exchange between stellar wind protons and the extended upper neutral atmosphere.

18 18 H. Lammer et al. Fig. 7 Preliminary model simulation of an expanded hydrogen-rich thermosphere of 7R Earth and related hydrogen corona which contains outward flowing planetary hydrogen atoms and stellar wind produced hydrogen ENAs around an Earth-like planet within an orbit of an M-star habitable zone at 0.16 AU. The dark circle correspond to the size of the planet. x-axis points towards the star. Right, illustration of possible size relations of terrestrial exoplanets with an extended hydrogen coronae around M-dwarfs (after Lammer, et al. 2011b). 7 Conclusion Studies related to the origin of protoatmospheres of terrestrial planets indicate that dense water vapour dominated H 2 O/CO 2 -rich steam atmospheres with pressure levels from bars equivalent to an amount of hydrogen of Earth oceans can be outgassed during the magma ocean solidification process, additionally to possible accumulated nebular-based hydrogen envelopes which may remain from the early stage of protoplanet formation. If the atmospheric escape processes or the giant impact history cannot remove these initial hydrogen and water inventories a terrestrial planet can accumulate a dense abiotic oxygen-rich atmosphere or remains as a sub-uranus body even within the habitable zone of its host star. On the other hand if terrestrial planets would evolve during the high EUV-phase of their young host stars to N 2 -rich Earth-like atmospheres too early they may lose its nitrogen inventory. This indicates that Earth s N 2 -inventory should have been protected against escape by higher amounts of IR-cooling CO 2 contents in the thermosphere or dense hydrogen or oxygen envelopes. If the initial outgassed water inventory on early Earth would have been much larger than 500 bar, or the Moon-forming impact would not have removed a part of it, the Earth may have also accumulated an abiotic oxygen-rich atmosphere or evolved in a water world with no continents and a H 2 O/CO 2 atmosphere surrounded by an atomic H envelope. If the Earth would have originated dryer, the young Sun may have been active enough to remove the less dense initial atmosphere and the planet may have ended with a present day Venus- or Mars-like CO 2 atmosphere. Hence it follows that Earth-like planets within HZs of less active F-stars may have a problem to lose their initially dense outgassed protoatmospheres, while more