School of Engineering and Natural SciencesUniversity of Iceland

2013

School of Engineering and Natural SciencesUniversity of Iceland

2013

Planetary atmospheresLife in the Universe, EÐL 620M

Andri Geir JónassonÁrni JohnsenErica Massey

Sigurður ThorlaciusSnædís Björgvinsdóttir

Contents

List of Figures v

List of Tables vii

1 Introduction 1

2 Formation and evolution of atmospheres 32.1 Formation of atmospheres . . . . . . . . . . . . . . . . . . . . . . . . 32.2 Evolution of atmospheres . . . . . . . . . . . . . . . . . . . . . . . . . 52.3 Primary and secondary atmospheres . . . . . . . . . . . . . . . . . . 10

3 The atmosphere of Earth 133.1 Formation and evolution of Earth’s atmosphere . . . . . . . . . . . . 133.2 The Earth’s atmosphere today . . . . . . . . . . . . . . . . . . . . . . 163.3 Stratification . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 173.4 Biogeochemical cycles . . . . . . . . . . . . . . . . . . . . . . . . . . . 19

4 Atmospheres in our solar system 294.1 Atmospheres of planets . . . . . . . . . . . . . . . . . . . . . . . . . . 294.2 Atmospheres of moons . . . . . . . . . . . . . . . . . . . . . . . . . . 32

5 Atmospheres outside our solar system 355.1 The habitable zone . . . . . . . . . . . . . . . . . . . . . . . . . . . . 355.2 Exoplanets . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 365.3 Detection of exoplanetary atmospheres . . . . . . . . . . . . . . . . . 38

6 Terraforming Mars 436.1 Main problems . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 436.2 Advantages . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 446.3 Terraformation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 456.4 Comet impact . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 47

7 Conclusion 49

References 51

iii

List of Figures

2.1 Volcanic outgassing . . . . . . . . . . . . . . . . . . . . . . . . . . . . 5

2.2 Maxwell distribution . . . . . . . . . . . . . . . . . . . . . . . . . . . 7

2.3 Thermal escape . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 8

2.4 Charge exchange . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 9

2.5 Escape velocity and surface temperature . . . . . . . . . . . . . . . . 11

3.1 Evolution of Earth’s atmosphere . . . . . . . . . . . . . . . . . . . . . 14

3.2 Volcanic gas . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15

3.3 Composition of Earth’s atmosphere . . . . . . . . . . . . . . . . . . . 16

3.4 Gradients in Earth’s atmosphere . . . . . . . . . . . . . . . . . . . . . 18

3.5 The water cycle . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21

3.6 The carbon cycle . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 23

3.7 The nitrogen cycle . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24

3.8 The oxygen cycle . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 26

3.9 The sulfur cycle . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27

4.1 Rocky planets . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 30

4.2 Gas planets . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 31

4.3 Escape velocity . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 31

v

LIST OF FIGURES

4.4 Titan’s atmosphere . . . . . . . . . . . . . . . . . . . . . . . . . . . . 32

4.5 Jupiter’s moons . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 33

5.1 Habitable zone . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 36

5.2 Transit . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 37

5.3 Bohr atom . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 40

5.4 IR spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 41

6.1 The Martian atmosphere . . . . . . . . . . . . . . . . . . . . . . . . . 45

6.2 Dynamics of the Martian atmosphere . . . . . . . . . . . . . . . . . . 46

6.3 Comet orbit . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 48

vi

List of Tables

3.1 Composition of Earth’s atmosphere . . . . . . . . . . . . . . . . . . . 17

3.2 Reservoirs of water and their residence times . . . . . . . . . . . . . . 22

3.3 Reservoirs of carbon and their residence times . . . . . . . . . . . . . 22

3.4 Reservoirs of nitrogen and their residence times . . . . . . . . . . . . 25

3.5 Reservoirs of oxygen and their residence times . . . . . . . . . . . . . 25

3.6 Reservoirs of sulfur and their residence times . . . . . . . . . . . . . . 25

vii

1 Introduction

An atmosphere can be defined as a layer of gases gravitationally bound to a planet.The Earth’s atmosphere was initially formed 4.6 billion years ago, when our solarsystem was born and it is the only atmosphere known to be able to support life. Un-derstanding the formation, evolution and properties of atmospheres is essential whenlooking for life elsewhere in the universe. Considering that, it seems quite straight-forward that methods were developed to detect the composition of atmospheres onother celestial bodies. If a life-harboring atmosphere was found on another planet,it would open up new depths of worlds and might even reveal how life was formed.

Here is presented a brief summary of the main formation and evolutionary processesof planetary atmospheres. Furthermore, the atmosphere of Earth and other bodiesin the solar system are analyzed in detail and observational results of atmospheresoutside our solar system are summarized. Lastly, the possibility of creating anartificial atmosphere on Mars is explored.

1

2 Formation and evolution ofatmospheres

An atmosphere can be defined as a gravitationally bound reservoir of species volatileenough to exist as gases above a planetary surface, and dense enough so that itsphysics are dominated by collisional interactions. This definition of an atmospherewill be used throughout this section, and concludes that out of the terrestrial planets,Venus, Mars and Earth have these kind of atmospheres but Mercury does not.The space above Mercury (similarly, the Moon) is not empty but is consideredcollisionless, as the population of atoms is very thin. The “atmosphere” is episodic,transient and does not preserve clues to evolutionary history (Pepin, 2006). Inorder to fully describe the functions and properties of atmospheres, the formationand evolution of atmospheres with time must first be understood.

2.1 Formation of atmospheres

Understanding how the so-called primordial atmospheres evolved to their presentstate is closely related to the challenge of knowing the evolutionary path of theplanets themselves. The sun, the planets, and their atmospheres are believed tohave condensed about 4.6 billion years ago from a primitive solar nebula. Thecomposition of this nebula is believed to have been similar to the composition of thesun, mostly hydrogen and helium, but also a small amount of the heavier elements.The heavier elements in the nebula formed cosmic dust, which stuck together andformed small bodies called planetesimals. These planetesimals then accreted over aperiod of 0.1 to 1 million years and came to be larger bodies which merged over a100 million year period and formed the planets (Hunten, 1993; Pepin, 2006).

As the planets grew they had several opportunities to acquire the gaseous cloakcalled an atmosphere. The giant planets, Jupiter, Saturn, Uranus and Neptune, grewmassive enough to capture large quantities of hydrogen and other gases directly fromthe nebula while they were forming (Pollack et al., 1996). The terrestrial planets aremuch smaller and may have formed too slowly to gravitationally attract gases in thesame way the gas giants did. In addition, they are situated much closer to the Sun

3

2 Formation and evolution of atmospheres

in a region too hot for volatiles to exist as solid materials. There are several theorieson how the terrestrial planets were able to form atmospheres, three of which will becovered in the following chapters.

2.1.1 Gravitational capture

As mentioned before, the giant planets are believed to have acquired their atmo-spheres mainly by gravitational capture. For a terrestrial planet, however, the possi-bility of atmosphere formation happening by direct capture from the nebula requiresthe planets to have grown to a significant fraction of their present masses before thesolar nebula dissipated. The gravity of the planet would then have been strongenough to be able to pull in gases from the surrounding accretion disk. Accord-ing to observational and theoretical estimates this could be the case for Earth andVenus, the nebula could have survived long enough for these planets to capturesubstantial atmospheres. The mass of Mars is however ten times smaller than theEarth’s, so it might not have been massive enough to have condensed ambient gasesby gravitational capture (Pepin, 2006).

2.1.2 Cometary impacts

Another option for the formation of atmospheres is that volatiles could have beencarried into the growing planets by infalling planetesimals and dust. Comets, forexample, are made mostly of frozen volatile materials, such as water, carbon dioxide,methane and ammonia. If enough cometary impacts occurred, substantial amountsof these volatiles could have been carried by the comets to the planetary surfacesand captured by gravity. These impacts could also have eroded parts of an existingatmosphere of a planet. The balance is dependant on the composition and mass ofthe impactor and the mass of the growing planet. As an example, erosion on Earthand Venus became less efficient when the planets had grown to nearly their presentmasses. Mars, on the other hand, remained rather vulnerable and would mostlikely still be affected if the population of impactors had not decreased considerably.It remains difficult to know whether enough planetesimals may have run into theterrestrial planets to account for their atmospheres (Hunten, 1993).

2.1.3 Volcanic outgassing

The third possibility of how the terrestrial planets gained their atmospheres is vol-canic outgassing, shown on figure 2.1. This theory assumes that the early planets

4

2.2 Evolution of atmospheres

lost hydrogen, helium and other gases resulting in planets without atmospheres.Gases trapped in the interiors of planets would then have contributed to exterioratmospheres through a process called volcanic outgassing, where huge amounts ofgases are released during volcanic eruptions (Pepin, 2006).

Figure 2.1: Volcanic outgassing; gases previously trapped in the mantle are re-leased in during volcanic eruptions.

2.2 Evolution of atmospheres

2.2.1 Evolutionary processes

The previous chapter discussed possible ways for planets to gain atmospheres, in-troducing processes that involved atoms or molecules getting bound to a planet bygravity and creating a reservoir called an atmosphere. If atmospheres are to evolve,there must also be a way for atoms and molecules to leave the system. In these mat-ters, escape velocity is important because a substance has to overcome the escapevelocity in order to leave a planet.

Escape velocity is the minimum initial velocity required for a projectile fired verti-cally at a planet’s surface to escape its gravitational force. Skipping the details onhow it can be derived, escape velocity, v0, of a planet can be calculated as follows,where g is the gravitational acceleration, R is the radius of a planet and M is itsmass: v0 =

√2gM/R. As an example, Earth has the escape velocity of 11.2 km/s

5

2 Formation and evolution of atmospheres

but Jupiter’s escape velocity is 59.5 km/s. This means that it is much more difficultfor atoms or molecules to leave Jupiter than Earth (Sternheim & Kane, 1991).

There are several ways in which atoms and molecules can reach escape velocity,covered in the following chapters. In thermal escape, gases get too hot for gravityto hold on to them, while in non-thermal processes, chemical or charged-particlereactions throw atoms and molecules out of the atmosphere. In a third process,impact erosion, asteroid and comet impacts blast away the air (Catling & Zahnle,2009).

2.2.2 Thermal escape

The average velocity of the molecules in some volume of gas is determined by tem-perature, but in reality the velocities of individual molecules are always changing.Molecules in the gaseous phase are constantly colliding with each other, gaining andlosing kinetic energy in the process. Kinetic energy is determined by the mass of amolecule and its velocity by the following equation

Ekin =1

2mv2

Since the mass of a molecule is constant it can be concluded that the velocities ofindividual molecules should vary continuously with changes in kinetic energy. TheMaxwell distribution shown in figure 2.2 describes the variation in kinetic energyamong individual molecules. When individual molecules in the high velocity tailof the Maxwell distribution reach escape velocity and leave the atmosphere it iscalled thermal escape, more specifically Jeans escape. This can be described as airevaporating as atoms or small particles off the top of the atmosphere, above the so-called exobase. At such a high altitude the air is tenuous and the gas particles hardlyever collide. This means that there is nothing stopping the atoms or molecules withsufficient velocity from escaping into space (Catling & Zahnle, 2009). Lighter atomsare more susceptible to this mechanism than the more massive ones, because heavieratoms have lower average velocity of molecules at a given temperature. Jeans escapeis therefore thought to account for the deficiency of light atoms such as H and He inthe atmospheres of the terrestrial planets, as the lightest gases overcome a planet’sgravity most easily after having reached the exobase. In addition, Jeans escape canalso partially explain why the most massive bodies in the solar system have denseatmospheres (Hunten, 1993).

There is another way in which thermal escape can occur, called hydrodynamic es-cape. It is far more dramatic than Jeans escape, where gas evaporates at a molecularlevel. In hydrodynamic escape, also known as planetary wind, heated air escapesa planet’s gravity as a bulk mass. Moving a large mass of atoms requires a source

6

2.2 Evolution of atmospheres

Figure 2.2: The Maxwell distribution.

of energy, either by heat from the Sun or earlier in the evolutional history of at-mospheres from planetary accretion processes. The upper part of an atmospherecan absorb radiation from the sun, which causes it to warm and expand as the gasparticles gain more energy. The air mass rises smoothly and gains enough velocityto escape the planet. Hydrogen rich atmospheres are most susceptible to hydrody-namic escape, and as the air mass flows outwards it can pick up and drag alongheavier atoms and molecules (Catling & Zahnle, 2009).

Evidence of thermal escape is derived from which planets have atmospheres andwhich do not. This is shown in figure 2.3, where the strength of stellar heatingrelative to the strength of gravity seems to be the deciding factor. Planets withoutatmospheres have have strong heating and weak gravity but bodies with atmosphereshave weak heating and strong gravity (Catling & Zahnle, 2009).

2.2.3 Non-thermal escape

Non-thermal escape is the outcome of a single event where chemical reactions orparticle collisions cause atoms to gain escape velocity. These events, like the thermalescape events, take place above the exobase, where collisions are infrequent and donot hinder the escaping particle. Non-thermal processes commonly involve ions asthe escapees, charged particles usually bound to a planet by its magnetic field. Non-thermal escape is more important than thermal escape on some planets, includingthe Earth (Catling & Zahnle, 2009).

7

2 Formation and evolution of atmospheres

Figure 2.3: Evidence of thermal escape. Airless bodies have strong heating andweak gravity but bodies with atmospheres have weak heating and strong gravity(Catling & Zahnle, 2009).

Non-thermal escape is perhaps best explained with examples of such events. Onetype of event, known as charge-exchange, involves a fast hydrogen ion colliding witha hydrogen atom and capturing its electron. Only a small part of the energy ofthe fast ion is transferred to the electron donor, while most of it is retained by theresulting neutral atom. Since it retains most of its energy, it is still moving fast,and because it is not charged anymore it is immune to a magnetic field. In otherwords, the ion and the hydrogen atom have exchanged identities but not energies,making escape possible. In another process, polar wind, ions are able to escape.Most magnetic field lines reach from one magnetic pole to the other, but some getdragged out of place by the solar wind and are not able to move back. This resultsin a weak spot in the planet’s magnetic field where ions can escape. These ionsstill need to reach escape velocity, but are not bound by the magnetic field in thesespecific places so they drift off into space as a stream of charged particles.

Yet another non-thermal process is called photochemical escape. Radiation from theSun ionizes oxygen, nitrogen and carbon monoxide molecules which drift into the

8

2.2 Evolution of atmospheres

Figure 2.4: Charge exchange, an example of non-thermal escape.

upper atmosphere. Once the molecules have been ionized they either recombine withelectrons or collide with another charged molecule. The collision releases energy andsplits the molecules into atoms which are capable of escaping. This process is, forexample, known to be operating on Mars (Catling & Zahnle, 2009).

Planets without a global magnetic field are defenseless against solar winds, allowingsputtering, the last process covered in this chapter. The magnetic solar wind is ableto pick up ions, which then undergo charge exchange and escape into space. Mars,Titan and Venus, for example, do not have global magnetic fields and are thereforevulnerable to sputtering. This might then explain why the atmosphere on Mars isrich in heavy nitrogen and carbon isotopes (Catling & Zahnle, 2009).

2.2.4 Impact erosion

A far more drastic escape than occurs during thermal and non-thermal escape iswhen comets or asteroids crash into planets. Depending on the composition ofthe impactor and the mass of the planet, the impact can either add volatiles tothe existing atmosphere, as previously explained, or erode parts of it. When thecolliding objects are travelling at high speed and are sufficiently big, they vaporizeboth themselves and a similar mass of the surface at impact. Because the erodinggases expand faster than escape velocity, they are able to expel the overlying air intospace. Impact erosion has the most dramatic effects when the planet has a weakgravity and the colliding object has a high speed (Catling & Zahnle, 2009; Hunten,1993).

9

2 Formation and evolution of atmospheres

2.3 Primary and secondary atmospheres

2.3.1 Primary atmospheres

Atmospheres are often divided into two categories, primary and secondary atmo-spheres. Primary atmospheres are mostly composed of hydrogen (90+%) and helium(5+%) that accreted during planet formation. These light elements are then lostdue to the escape velocity of a planet, the element’s atomic mass, the temperatureat the surface or due to other factors. Escape velocity is linearly related to a planet’smass and radius (Schombert, 2013).

Atoms that move slower than the escape velocity of the planet stay in the atmo-sphere. If atoms move faster than a planet’s escape velocity, they are forced outsidethe gravitational field into space. A higher temperature and a lower molecular masswill result in a higher mean velocity of molecules so these factors directly affect thecomposition of an atmosphere. Planets closer to the Sun are warmer and thereforea single element will have a higher velocity on a planet close to the Sun rather thana planet distant from the Sun (Schombert, 2013).

The relationship of escape velocity and surface temperature is illustrated in figure2.5.

2.3.2 Secondary atmospheres

The four planets closer to the Sun, specifically Mercury, Mars, Earth and Venus,have lost most of their lighter elements such as hydrogen and helium. As a result,they form a secondary atmosphere from tectonic activity and outgassing. The ma-jority of elements are rocky (iron, olivine and pyroxene) or volatiles classified asicy materials (water, carbon dioxide, methane, ammonia and sulphur dioxide) thatmix in a planet’s early mantle and crust. Icy materials are volatiles with extremelylow melting points but over 100 K (Volatiles , 2013). At this stage, icy materialsevolve into gases that move up to the surface by outgassing of water, carbon dioxide,sodium dioxide, nitrogen and noble gases, to form a secondary atmosphere. Sur-face temperature and interior chemistry of the planet determine outgassing (i.e. icymaterials can stay within the mantle if a planet cools quickly resulting in little tono tectonic movement). On the contrary, significant volcanic and tectonic activityoccurs if heat is kept within the mantle during accretion, creating planets with alarger mass (Schombert, 2013).

10

2.3 Primary and secondary atmospheres

Figure 2.5: Relationship between surface temperature (km/s) and escape velocity(K) of planets and moons. A gas escapes from a celestial body if it’s below thecorresponding line (Rieke, 2013).

11

3 The atmosphere of Earth

The modern Earth has a secondary atmosphere and it is (not surprisingly) the bestknown atmosphere and also the one that has been researched the most. In order tounderstand the properties and functions of atmospheres, it will now be described insome detail.

3.1 Formation and evolution of Earth’s atmosphere

In the following chapters, time will be represented with Ga (gigaannum, 109 years)which stands for billion years. Earth’s geological history is often divided into foureons: The Hadean eon from 4.57 to 3.8 Ga ago, the Archean eon from 3.8 to 2.5 Gaago, the Proterozoic eon from 2.5 to 0.54 Ga ago and the Phanerozoic eon from 0.54Ga ago until the current day (Zamora, 2013). The evolution of the atmosphere’scomposition over the eons is shown on figure 3.1.

3.1.1 Hadean eon

The Earth coalesced and formed 4.57 Ga ago and soon organized itself into layerswith heavier compounds at the core and lighter compounds at the surface. The coreof circulating molten metal established the magnetosphere, the magnetic field ofthe Earth that deflects most solar wind ions. The outermost layer was the originalprimary atmosphere. It was mostly composed of hydrogen and helium along withtrace compounds, similar to the solar nebula in which the solar system formed(Zamora, 2013).

These compounds escaped the Earth‘s atmosphere through thermal and non-thermalescape and what was left was mainly methane, ammonia and water. A large partof the Earth‘s water and nitrogen is thought to have come from comet impacts,probably in the Late Heavy Bombardment which took place in the Hadean eonaround 3.9 Ga ago. This bombardment of comets kept the surface molten (Zamora,2013).

13

3 The atmosphere of Earth

Figure 3.1: The evolution of the composition of the Earth’s atmosphere (Zamora,2013).

The high surface temperature drove the endothermic reaction of methane with watersteam shown in equation 3.1. The resulting carbon monoxide readily combines withmetals to form carbonyl compounds, therefore reducing the amount of carbon in theatmosphere as can be seen in figure 3.1 (Zamora, 2013).

CH4 +H2O→ CO+ 3H2 (3.1)

3.1.2 Archean eon

When the Late Heavy Bombardment was over and the crust had hardened, carbondioxide, water and other substances were emitted through volcanic outgassing anderuptions. The atoms of these volcanic gases had been bound in minerals when theEarth formed, but were released through eruptions and volcanic emissions (Braga-son, 2010). The main species of volcanic gases that are emitted through volcanicvents are shown in 3.2. Soon thereafter, the temperature was low enough for wa-ter to condense, forming oceans and dissolving carbon dioxide and ammonia. Thedissolved ammonia could form ammonium compounds, amines and other nitrogen-containing substances needed for the origin of life (Zamora, 2013). When carbon

14

3.1 Formation and evolution of Earth’s atmosphere

Figure 3.2: Average composition of volcanic gas at a volcanic vent (Textor et al.,2003).

dioxide dissolved in water, it reacted with calcium to form solid calcium carbon-ate that settled to the ocean floor, reducing the amount of carbon dioxide in theatmosphere (Bragason, 2010).

At that time, there was little to no oxygen and ozone in the atmosphere so ultravi-olet light could easily penetrate the atmosphere. The energy-rich light broke downthe ammonia in the atmosphere to nitrogen, which is a stable molecule and unreac-tive. The aerobically respiring cyanobacteria (blue-green algae) evolved during theArchean eon (Bragason, 2010).

3.1.3 Proterozoic eon

In the Proterozoic eon, the chemically inert nitrogen gas (N2) became the majorcomponent of the atmosphere once the other gases had escaped, precipitated asliquids or chemically reacted to form solid minerals. The cyanobacteria producedoxygen and carbohydrates with solar energy, carbon dioxide and water throughphotosynthesis, described by equation 3.2 (Zamora, 2013).

6CO2 + 6H2O+ energy→ C6H12O6 + 6O2 (3.2)

The oxygen produced by the cyanobacteria went into oxidizing metals and methaneon Earth. This can be seen in geological records as iron bands in sea layers andiron beds in land layers when the ocean had been saturated. After some time, mostof the metals had been oxidized and oxygen started building up in the atmosphere(Hreggviðsson, 2013).

15

3 The atmosphere of Earth

Some of the oxygen migrated to the upper atmosphere where it was transformedto ozone by ultraviolet light, forming the ozone layer. The formation is shown inequation 3.3 where M stands for an atmospheric compound like N2 or O2 (PressRelease: The 1995 Nobel Prize in Chemistry., 1995). The ozone layer hinders high-energy ultraviolet light from reaching the Earth‘s surface, making it possible foroxygen-respiring organisms to colonize the surface of the Earth (Zamora, 2013).

O2 + energy→ O+O O+O2 +M→ O3 +M (3.3)

In the Phanerozoic eon, the ratio of oxygen in the atmosphere fluctuated around 20%except for one catastrophic drop from 30% to 12%, causing a great mass extinction(Zamora, 2013).

Figure 3.3: The composition of the atmosphere shown on a logarithmic scale.Water vapor concentration is not included in the calculations. It can vary from0.001-5% and the average is about 1% (Atmosphere of the Earth, 2013; Williams,2010).

3.2 The Earth’s atmosphere today

The Earth is a dynamic system where geological and life processes interact andprovide the environment in which we live. During the last billion years in thehistory of the Earth, the average composition of the atmosphere, the oceans and theland has remained relatively constant (vanLoon & Duffy, 2011).

16

3.3 Stratification

The Earth’s atmosphere is a shell of gases surrounding the globe. Most humanactivity takes place in the lowest 20 km of the atmosphere. The relative ratios ofthe major gases remain relatively constant up to an altitude of 80 km (vanLoon &Duffy, 2011).

The substance composition of the Earth’s atmosphere is shown in table 3.1. Sincethe concentration of water vapor is highly variable, the relative ratios are calculatedwithout water. The concept of residence time is explained in chapter 3.4.

Table 3.1: The composition of the atmosphere (Atmosphere of the Earth, 2013)and the average time a molecule resides in the atmosphere (Railsback, 2013). *Watervapor concentration is not included in the calculations. It can vary from 0.001 - 5%and the average is about 1% (Atmosphere of the Earth, 2013; Williams, 2010).

Gas Volume Residence time in[ppmv] [%] the atmosphere

N2 Nitrogen 780840 78.084 107 yearsO2 Oxygen 209460 20.946 3000 - 10000 years

H2O Water vapor* 10000 1 10 daysAr Argon 9340 0.934 Forever

CO2 Carbon dioxide 394.45 0.039445 2 - 10 yearsNe Neon 18.18 0.001818 ForeverHe Helium 5.24 0.000524 1000000 yearsCH4 Methane 1.79 0.000179 2 - 10 yearsKr Krypton 1.14 0.000114 ForeverH2 Hydrogen 0.55 0.000055 4 - 8 years

N2O Nitrous oxide 0.325 0.0000325 5 - 200 yearsCO Carbon monoxide 0.1 0.00001 60 - 200 daysXe Xenon 0.09 0.000009 ForeverO3 Ozone 0 - 0.07 0 - 0.000007

NO2 Nitrogen dioxide 0.02 0.000002I2 Iodine 0.01 0.000001

3.3 Stratification

The atmosphere is divided into four sections based on the temperature gradient. Thegradient is shown on figure 3.4 and different layers are explained in correspondingsections: the troposphere, the stratosphere, the mesosphere and the thermosphere.Note that thermodynamic temperature is shown, which is the measure of the kineticenergy of molecules. Because there are few molecules in the thermosphere, a con-ventional thermometer would not give a reading as high as figure 3.4 shows vanLoon& Duffy (2011).

17

3 The atmosphere of Earth

(a) Temperature (b) Pressure

Figure 3.4: Temperature and pressure gradients of Earth’s atmosphere as a func-tion of altitude.

3.3.1 Troposphere

The temporal and spatial average temperature at the Earth’s surface is 14◦ C (287K). From the surface and upwards, the temperature decreases until the tropopause(the boundary between the troposphere and the stratosphere) is reached at an al-titude of about 15 km. The troposphere contains approximately 85% of the atmo-sphere’s mass (vanLoon & Duffy, 2011) and nearly all of the atmosphere’s watervapor and aerosols. It is thoroughly mixed because the air at lower altitudes iswarmer and tends to rise (Russell, 2010).

3.3.2 Stratosphere

In the stratosphere, the temperature increases with altitude. Such a temperatureprofile creates very stable atmospheric conditions, so little mixing takes place. Itreaches up to the stratopause at approximately 50 km (vanLoon & Duffy, 2011).

The stratosphere differs from the troposphere in ozone concentration. Due to in-creased solar radiation at higher altitudes, ozone is synthesized and an ozone layeris formed (vanLoon & Duffy, 2011). Further explanation of the ozone layer is inchapter 3.1.3.

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3.4 Biogeochemical cycles

3.3.3 Mesosphere

In the mesosphere, temperature decreases with altitude. The mesosphere is wheremost meteors burn up and it reaches up to the mesopause at an altitude of approx-imately 85 km (Russell, 2008).

3.3.4 Thermosphere

In the thermosphere, temperature increases with altitude. Charged particles (elec-trons, protons, and other ions) from space collide with molecules in the thermo-sphere, exciting them into higher energy states. These molecules emit energy asphotons, producing aurora (the Northern Lights).

3.3.5 Exosphere

The division between the thermosphere and the exosphere is called the thermopauseor the exobase which is located at an altitude of 500 km to 1000 km (Russell, 2008).Light gases above the exobase can overcome the Earth’s gravity and escape theatmosphere through thermal or non-thermal escape.

3.3.6 Ionosphere

The ionosphere forms the inner edge of the magnetosphere and covers the ther-mosphere and a part of the exosphere. It is a shell of electrons and electricallycharged atoms and molecules that surrounds the Earth. It exists mainly due to theultraviolet radiation from the Sun (Ionosphere, 2013).

3.4 Biogeochemical cycles

A description of the composition of the Earth’s environment is not exhaustive be-cause the system is dynamic and there are many physical, chemical and biologicalprocesses that operate within and link components of the environment. In order todescribe this dynamic system, it is often divided into four principal compartments:the atmosphere (the gaseous environment), the hydrosphere (the liquid environ-ment), the lithosphere (the solid, terrestrial environment) and the biosphere (the

19

3 The atmosphere of Earth

living environment). Chemical processes are described within each environmentand the reactions that transmit from one to the other (vanLoon & Duffy, 2011).

Most of the elements on Earth required for life are continually cycling through theEarth’s principal compartments on time scales from a few days to millions of years.These cycles are called biogeochemical cycles (Biogeochemical Cycles , 2008).

Many elements go through such cycles and some of them are essential to life onEarth, including carbon, nitrogen, oxygen, hydrogen, sulphur and phosphorus. Inthat sense, the elements are recycled but the rate of substance cycling can varygreatly and in some places the accumulated element can be held for a very longtime, even for millions of years (Biogeochemical Cycles , 2008).

In a steady state, the parts of a cycle are balanced so that concentrations remainconstant and the average time an element stays in each part is called the residencetime (vanLoon & Duffy, 2011).

Many of the cycles are driven by organisms, such as nitrogen-fixing bacteria in thenitrogen cycle. The cycles could therefore be signs of life in the search for life onother planets.

3.4.1 Water cycle

Water is continuously cycled between its various reservoirs. The cycling can occurby many processes like evaporation, condensation, precipitation, runoff, infiltration,sublimation, transpiration, melting, and groundwater flow (Pidwirny, 2006). Oneexample of a sub-cycle is the water in the atmosphere that precipitates to the surface,infiltrates to the lithosphere, flows as groundwater to the ocean (hydrosphere), whereit evaporates to the atmosphere once again. Figure 3.5 shows a simplified version ofthe water cycle.

The volume of the main reservoirs of water is shown on table 3.2 along with theaverage residence time for each reservoir.

3.4.2 Carbon cycle

The carbon cycle is often divided into the short-term and the long-term carboncycle. In the long-term carbon cycle, CO2 is emitted from the lithosphere to theatmosphere through volcanic outgassing. The CO2 in the atmosphere enters thehydrosphere by dissolving in water and forming HCO3. It is then transformed from

20

3.4 Biogeochemical cycles

Figure 3.5: The main reservoirs and fluxes in the water cycle.

the hydrosphere back to the lithosphere when CaCO3 settles and gets buried insettlements, eventually becoming sedimentary rock (Gíslason, 2012).

In the short-term carbon cycle, CO2 flows from the atmosphere to the biospherethrough photosynthesis and goes back through respiration. A similar sub-cycle iswhen CO2 of the atmosphere gets dissolved in water and evaporates again (Gíslason,2012). Figure 3.6 shows the short-term carbon cycle.

The volume of the main reservoirs of carbon is shown on table 3.3 along with theaverage residence times for each reservoir (Gíslason, 2012). It is also shows whichcarbon cycle (the long-term or the short-term carbon cycle) the reservoir mainlybelongs to. Note that the numbers in table 3.3 and figure 3.6 do not entirely matchdue to uncertainty and anthropogenic effects. The 6.4 GtC released by the burningof fossil fuels shown in figure 3.6 is for the year 1994 but the carbon release was 8.7GtC in 2008 (Gíslason, 2012). The carbon cycle would be in a steady state if notfor this anthropogenic effect. The carbon in fossil fuels is taken from the long-termcarbon cycle and thrown into the short-term cycle every year by humans throughburning, thus disrupting both cycles. The effect is more drastic in the short-termcarbon cycle, since the reservoirs are smaller. The CO2 released is a greenhouse gas

21

3 The atmosphere of Earth

Table 3.2: Reservoirs of water and their residence times (Pidwirny, 2006). *Theresidence time of ice in Antartica is very long, approximately 20000 years (Watercycle, 2013).

Reservoir Volume Fraction Residence time[Mkm3] [%]

Oceans 1370 97.25 3200 yearsIce Caps and Glaciers* 29 2.06 20 to 100 yearsGroundwater 9.5 0.67 100 to 10.000 yearsLakes 0.125 0.01 50 to 100 yearsSoil Moisture 0.065 0.005 1 to 2 monthsAtmosphere 0.013 0.0009 9 daysStreams and Rivers 0.0017 0.0001 2 to 6 monthsBiosphere 0.0006 0.00004Total 1408.7

Table 3.3: Reservoirs of carbon and their residence times (Gíslason, 2012).Reservoir Carbon cycle Size Residence time

[GtC] [years]Sedimentary rock Long 5 · 107 2 · 108Sea Short/Long 39000 360Fossil fuels Long 4200 540Soil Short 1600 20Atmosphere Short/Long 770 3.6Organisms Short 610 15

which leads to global warming.

3.4.3 Nitrogen cycle

Nitrogen is essential for living organisms to produce complex organic molecules likeamino acids, proteins, and nucleic acids. There is an abundance of nitrogen in theatmosphere (as N2) as can be seen on figure 3.7 but it is chemically nonreactive andmost plants can only take up nitrogen in two solid forms: ammonium ion (NH+

4 )and the ion nitrate (N–

03) (Pidwirny, 2006).

Almost all of the nitrogen found in terrestrial ecosystems originally came from theatmosphere. Some was transported to the biosphere through rainfall or the effectsof lightning but the majority came from biochemical fixation by nitrogen fixingbacteria (Pidwirny, 2006).

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3.4 Biogeochemical cycles

Figure 3.6: The main reservoirs and fluxes of the short term-carbon cycle (Solomonet al., 2007).

When organisms decay, nitrogen is converted into inorganic forms through decom-position. Decomposers (bacteria and fungi) in the upper soil layer chemically mod-ify the nitrogen found in organic matter from ammonia (NH3) to ammonium salts(NH+

4 ) through mineralization (Pidwirny, 2006).

Most of the ammonium is chemically altered by bacteria into nitrite (NO–2) and

then to nitrate (NO–3) by another bacteria through nitrification. The nitrate is

very soluble and leaches to the oceans where it can be returned to the atmospherethrough denitrification. The process of denitrification involves yet another bacteriathat reduces nitrate (NO–

3) into either nitrogen (N2) or nitrous oxide (N2O) gas thatdiffuses into the atmosphere (Pidwirny, 2006). The nitrogen cycle is shown on figure3.7.

3.4.4 Oxygen cycle

Oxygen was not found in the atmosphere until stromatolites (biofilms of microor-ganisms) and photosynthetic cyanobacteria evolved during the Archean eon as wasexplained in chapter 3.1.2. The main driver of the oxygen cycle is photosynthesis,where oxygen and carbohydrates are formed from carbon dioxide and water. The

23

3 The atmosphere of Earth

Figure 3.7: The main reservoirs and fluxes of nitrogen cycle.

reverse happens in respiration, where energy rich organic matter is used as an energysource, and decay, where organic matter is broken down. A simplified oxygen cycleis shown on figure 3.8, note that the flow of oxygen through the hydrosphere is notshown.

The largest reservoir of oxygen is in the lithosphere, in sedimentary rocks, as can beseen in table 3.5.

3.4.5 Sulfur cycle

Sulfur in its natural form is a solid and is transported by physical processes likewind, erosion by water, and geological events like volcanic eruptions. It can enterthe atmosphere, the biosphere and the hydrosphere in its compounds, such as sulfurdioxide, sulfuric acid and salts of sulfate or organic sulfur (Moses & Pidwirny, 2010).Volcanic eruptions emit sulfur directly into the atmosphere where it is converted intosulfate (SO4). The sulfate is taken up by plants and microorganisms and converted

24

3.4 Biogeochemical cycles

Table 3.4: Reservoirs of nitrogen and their residence times (Reeburgh, n.d.).Reservoir Size Residence time

[GtN] [years]Atmosphere N2 4 · 106 107

Sediments 5 · 105 107

Ocean 22000 1000Soils 95 2000Terrestrial biomass 35 50Atmosphere N2O 14 100Marine biomass 0.5 0.5

Table 3.5: Reservoirs of oxygen and their residence times (Reeburgh, n.d.).Reservoir Size Residence time

[GtO2]Sedimentary rocks 3.2 · 107Atmosphere 1.184 · 106 3 · 106 yearsOcean 7200 20 days - 500 yearsBiota 6112 50 - 1000 years

into organic forms. The sulfur is again released as a sulfate when the organisms dieand decay. The sulfur enters the hydrosphere as well and a portion of this sulfur isemitted back into the atmosphere from sea spray. The remaining sulfur is lost tothe ocean depths, combining with iron and entering the lithosphere (Sulfur Cycle,2006).

Table 3.6: Reservoirs of sulfur and their residence times (Reeburgh, n.d.).Reservoir Size Residence time

[GtS]Lithosphere 2 · 107 108 yearsOcean 3 · 106 106 yearsSediments 3 · 105 106 yearsSoils 300 1000 yearsLakes 0.3 3 yearsMarine Biota 0.03 1 yearAtmosphere 0.0048 8 - 25 days

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3 The atmosphere of Earth

Figure 3.8: The main reservoirs and fluxes of a simplified oxygen cycle. The largestreservoir of oxygen is in the lithosphere, in sedimentary rocks, as seen in table 3.5.

3.4.6 Other biogeochemical cycles

Other noteworthy biogeochemical cycles are the phosphorus cycle, the silica cycleand cycles of iron and other minerals. The phosphorus cycle is important for lifebut it does not take place in the atmosphere.

26

3.4 Biogeochemical cycles

Figure 3.9: The main reservoirs and fluxes of the sulfur cycle (Sulfur Cycle, 2006).

27

4 Atmospheres in our solar system

The previous section described the evolution of the Earth’s atmosphere and it’sproperties. Other celestial bodies can also have outgassing, volatiles, rocky materialsand atmospheres. Our solar system is relatively well known and the atmospheresof the planets will be described in section 4.1. The only planet in our solar systemthat lacks an atmosphere is Mercury. Earth’s moon has water-ice at its poles and noatmosphere or magnetic field, but that is not the case for all moons. Some moonsin our solar system have atmospheres which can be quite exotic and mysteriousatmospheres, these will be described in section 4.2.

4.1 Atmospheres of planets

There are eight planets in our Solar System. In order from closest to furthest fromthe Sun are Mercury, Venus, Earth, Mars, Jupiter, Saturn, Uranus and Neptune.Three of the four rocky planets closest to the Sun have a secondary atmosphere.Mercury, the rocky planet closest to the Sun, has an “atmosphere” that is essentiallya vacuum composed of 42% oxygen with less concentrations of hydrogen, nitrogenand potassium. Its temperature is highly variable, from −183◦ C to 427◦ C (Strobel,2013). The four planets further from the Sun have temperatures decreasing withdistance and primary atmospheres made of primarily hydrogen and helium. Incontrast to Mercury’s variable temperatures and lack of atmosphere, the secondplanet from the Sun, Venus, is the hottest planet because its atmosphere is thickwith greenhouse gases that trap and retain the Sun’s heat. Illustrated in figure 4.1,the atmospheres of Venus and Mars are composed similarly of 95% carbon dioxidewith variable other concentrations of nitrogen, oxygen, water, nitrogen oxide andsulphur dioxide unique to each planet. Earth’s atmosphere is composed of 78%nitrogen with less concentrations of oxygen, water and carbon dioxide. Figure 4.1compares atmospheric concentrations of Mercury, Venus, Earth and Mars.

Differences in atmospheric pressure and composition between Earth and Venus seemdifficult to understand. The two planets are similar in mass and distance from theSun, so it would be natural to assume they would have acquired similar primordialatmospheres in their youth and followed similar evolutionary paths to the present.

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4 Atmospheres in our solar system

Figure 4.1: The Solar System’s planets closest to the Sun have variability inatmospheric composition. A logarithmic scale illustrating Mercury having mostlyoxygen, Earth mostly nitrogen, Venus and Mars mostly carbon dioxide (Strobel,2013).

The volatile inventories of Venus and Earth are actually comparable if the carbondioxide locked up as carbonates in the Earth’s sedimentary column are included. Ifthis carbon dioxide were released in its entirety, the Earth would be transformedto a Venus-like planet with similar surface pressures. The reason for the differenceis likely caused by the difference in surface temperature that allows liquid water toexist on Earth, which takes up atmospheric CO2 and precipitates it as carbonaterock (Pepin, 2006).

The gaseous planets are composed of approximately 90% hydrogen and 10% helium,all with clouds of ammonia and water ice, with Neptune additionally having methaneice in its atmosphere. Figure 4.2 illustrates their atmospheric composition.

Temperature and atmospheric pressure relative to the Earth vary, shown in figure4.3, with Mercury and Mars having little to no pressure, to Venus having 92 timesthe pressure and gaseous planets having much greater than 1000 times the pressureof Earth. Venus and Mercury are the warmest planets with planets progressivelycolder at further distances from the Sun (Strobel, 2013).

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4.1 Atmospheres of planets

Figure 4.2: The Solar System’s four outer gaseous planetary atmospheres are com-posed primarily of hydrogen and helium. A logarithmic-scale comparison (Strobel,2013).

Figure 4.3: Temperature and pressure of planets in the Solar System relative toEarth (Strobel, 2013).

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4 Atmospheres in our solar system

4.2 Atmospheres of moons

4.2.1 Titan

Titan, Saturn’s largest moon, is the only natural satellite known to have a denseatmosphere, composed of 98% Nitrogen and also methane, shown in figure 4.4. Ithas a super-rotator atmosphere that rotates faster than its surface (Hecht, 2013).Titan is the only object other than Earth with stable bodies of surface liquid, such asorganic nitrogen-rich smog and clouds of methane and ethane at very cold temper-atures forming liquid methane rain (Titan: Moon, 2013). Titan’s methane cycle iscomparable to Earth’s water cycle. Titan has a subsurface ocean, lakes, hills, caves,sand dunes, fog, mist, smoggy haze, tropical cyclones and rain clouds, thought tolikely be a biotic environment able to support microbial extraterrestrial life (Hecht,2013). A mysterious fluorescent glow surrounds Titan, yet to be identified by scien-tists (Perkins, 2013).

Figure 4.4: Titan’s theorized internal structure (Model of Titan’s atmosphere,2013).

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4.2 Atmospheres of moons

4.2.2 Triton

Triton, Neptune’s moon, is the coldest body in the solar system with an extremelythin atmosphere composed of mostly nitrogen and methane. Due to Triton’s highalbedo, the mean annual surface temperature stays extremely cold at −235◦ CKnight (2013). Nitrogen-ice particles form thin clouds above the surface, boast-ing nitrogen and methane ice volcanoes.

4.2.3 Jupiter’s moons

Four of Jupiter’s sixty-three moons, the Galilean moons, were the first observed inthe 1600’s. In progression from closest to furthest from Jupiter they are namedIo, Europa, Ganymede and Calipso. Io is the most volcanically active body in thesolar system, spewing sulfurous plumes up to 500 km high (Erickson, 2010). Itsatmosphere is thin with a sodium atom halo cloud scattering sunlight at the yellowwavelength. Io’s atmosphere has energetic sulfur and oxygen ions that meet Jupiter’smagnetic field creating auroras, indicating that it likely has an ionosphere (Erickson,2012). Europa has a water ice surface covering an ocean of water and slushy ice,shown in figure 4.5, and believed to have twice the water as Earth. Ganymede is thelargest moon in the solar system and the only moon known to have an internallygenerated magnetic field.

Figure 4.5: Four of Jupiter’s moons: Io, Europa, Ganymede and Callisto. Theinternal structure changes the further the distance from their planet, Jupiter (Rieke,2013).

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5 Atmospheres outside our solarsystem

The Universe is immensely big and it has been estimated that there are 6 · 1022stars in the Universe (Gott et al., 2003). Eight planets orbit the Sun and most ofthem have an atmosphere so it is highly likely that a star detected outside the solarsystem has planets or moons with atmospheres. The atmosphere and the presenceof liquid water on the surface determines whether the planet is habitable for life aswe know it.

5.1 The habitable zone

Life as we know it is dependant on the presence of water. Therefore a circumstellarhabitable zone is traditionally defined as the region in which a terrestrial-mass planetcould sustain liquid water on its surface. The atmosphere of the planet wouldideally contain CO2, H2O and N2 with sufficient atmospheric pressure and an idealtemperature for enable water to exist. The acceptable temperature range does nothave defined values, as it is a function of a star’s luminosity and the distance fromplanet to star. In simple terms, the habitable zone is an area around a star in whichlife as we know it could possibly exist (Prantzos, 2006).

One of the primary goals of radial velocity and transit surveys is to find an earth-like planet, with a mass 0.3-10 times the mass of Earth situated in this so-calledcircumstellar habitable zone, shown in figure 5.1. The boundaries of the zone arerepresented by loss of water and maximum greenhouse effect. A planet that is closerto a star than the inner edge of the habitable zone has water-loss, therefore it willnot be able to hold liquid water on it surface. The temperature of the planet wouldbe too high and any water present would be lost into space. Outside the outerboundary of the habitable zone, the maximum greenhouse, it is too cold for liquidwater to exist. The maximum greenhouse limit is defined by the greenhouse effectfailing, when carbon dioxide begins to condense out of the atmosphere, making thesurface too cold for liquid water (Kopparapu et al., 2013).

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5 Atmospheres outside our solar system

Figure 5.1: The newly redefined cloud-free habitable zone (flux) boundaries forstars with different temperatures. The boundaries are represented by the water-lossand maximum greenhouse limit. The less conservative definition is also shown, withrecent boundaries of Venus and early Mars (Kopparapu et al., 2013).

5.2 Exoplanets

An exoplanet is a planet gravitationally bound to a parent star other than theSun. As of March 2013, there are 861 exoplanets known to mankind, most ofwhich are gas giants with primary atmospheres. The closest known exoplanet isAlpha Centauri Bb, located 4.37 light-years from Earth (The Extrasolar PlanetsEncyclopaedia, 2013).

The search for exoplanets is an essential step when looking for life in the Universe ifwe expect to find life in similar conditions to our own, on a planet orbiting a parentstar inside its habitable zone.

5.2.1 Methods of detection

A number of methods are used to find exoplanets. The two proven most effectiveare the radial velocity method and the transit method.

Radial velocity utilizes spectroscopy to measure the motion of the parent star, whichcan be changed by orbiting exoplanets. Heavy exoplanets that are close to their

36

5.2 Exoplanets

parent star have a greater effect on the parent star’s motion, therefore this methodis biased to finding heavy exoplanets that commonly are gas giants. As of March2013, a total of 501 exoplanets have been detected using this method (The ExtrasolarPlanets Encyclopaedia, 2013).

The transit method, illustrated in figure 5.2, measures changes in visual brightnessof a parent star caused by an orbiting planet crossing in front of the star. If this istruly an orbiting exoplanet, the change in brightness will happen periodically with aperiod which matches the orbital period of the planet. From Kepler’s third law, onecan then determine the distance between the planet and parent star, which givesimportant information to whether the planet is in the star’s habitable zone. Onemajor drawback of this method is that the planetary orbit must be aligned in a veryspecific manner, so that it will cross our line of sight to the parent star (Wright &Gaudi, 2012).

The Kepler space observatory is our most potent tool to detect exoplanets withthis method. Kepler monitors changes in brightness of more than 100.000 starsand determines if there are any periodic changes in brightness. Kepler has found2740 candidates as possible extrasolar planets. Using statistical methods, analystshave deduced that there are approximately 100-400 billion exoplanets in our galaxy(Billions and billions of Planets , 2013).

Figure 5.2: A diagram explaining the transit methods of detecting exoplanets. Thegraph shows the planets brightness as a function of time. If this drop in brightness isseen at least three times, we can confirm that the star has a planet orbiting it(Wright& Gaudi, 2012).

In total, as of March 2013, 294 exoplanets have been detected with the transitmethod. This number is bound to increase rapidly in the coming months or years, asmore data from Kepler is analyzed. Other methods than the two already mentionedhave only detected 66 exoplanets (The Extrasolar Planets Encyclopaedia, 2013).

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5 Atmospheres outside our solar system

5.2.2 Classification of habitability

In order to classify which exoplanets are most likely to be habitable for life, astrobi-ologists have several different criteria, the most important being the Earth similarityindex (ESI), which is a measure of how similar the planet is to Earth in terms ofradius, density, escape velocity and surface temperature. Earth has by definition anESI value of 1 and our closest neighbors, Venus and Mars, have ESI values of 0.78and 0.64, respectively (The Habitable Exoplanet Catalog , 2013).

The most prominent candidate for an Earth-like, habitable planet is Gliese 667C c,which has an ESI value of 0.79 and orbits inside the habitable zone of its parent star.Comparing that to Venus’ ESI value, this may not seem like a good candidate, butastrobiologists believe that Gliese 667C c’s surface could be suitable for vegetation,while Venus’ surface definitely is not (The Habitable Exoplanet Catalog , 2013).

5.3 Detection of exoplanetary atmospheres

When looking at possible candidates of exoplanets which could host life, the plane-tary atmosphere can give important clues to whether the planet could or is hostinglife. For life as we know it to be possible, there has to be present a number ofchemical compounds on the planet, the most important being water. Carbon com-pounds, such as carbon dioxide, methane and more complicated carbon chains arealso important since carbon is the main building block in the organic chemistry weare familiar with on Earth. Even the presence of some chemicals which we wouldnot immediately connect with life, such as hydrogen sulfide, can show us whereorganisms using chemosynthesis could be present (Hreggviðsson, 2013).

5.3.1 Spectroscopic measurements

The main method of detecting planetary atmospheres use spectroscopy to determinethe chemical composition of the atmosphere, or look for clues as to which chemicalsare present in it. The main idea of spectroscopy is to either heat up or shine lighton some unknown material and study the outgoing light after its been broken up bycolors, for instance by going through a prism. If some colors in the outgoing lightare exceptionally bright then they are called emission lines, but if some colors aren’tpresent in the outgoing light then they’re called absorption lines. Together, theselines are called spectral lines (Haken & Wolf, 2005).

Each element in the periodic table has a unique “signature”, a set of spectral lines

38

5.3 Detection of exoplanetary atmospheres

which is different from all other elements. In hot conditions, such as on the surfaceof a star, the elemental spectral lines will show as emission lines. In cold conditions,such as on the surface of a planet, they will show as absorption lines. However,when searching for a certain chemical compound, we need to know more than justwhich elements are present. We need to also know which chemical bonds theyhave formed. For example, water and hydrogen peroxide are made up of the sameelements, hydrogen and oxygen, but have vastly different chemical properties. Whilewater is the most important chemical compound for life, hydrogen peroxide is usedto bleach paper, whiten teeth and for sterilization.

Thankfully, spectroscopy is also helpful in this context. Every chemical bond has itsown “fingerprint”, although it’s not as sharp or easy to see as the ones from singleelements. These spectral lines are in the infrared part of the electromagnetic spec-trum, as opposed to the visible and ultraviolet spectral lines of individual elements.To explain why this is and from where spectral lines originate in general, we musttake a closer look at the building blocks of matter around us (Ágúst Kvaran, 2013).

The world of small things, like atoms and chemical compounds, is governed byquantum mechanics. Quantum mechanics predict that particles are not hard spheresbut travel as waves. The lighter the particle is, the more obvious its wave propertiesare, which is why we don’t experience this effect in the macroscopic world we live in.The particle’s wavelength is related to how much energy the particle has (Cohen-Tannoudji et al., 1977).

Now consider a single electron. If the electron is free to travel in any direction, it canhave any possible wavelength and thus any possible energy, so nothing interestinghappens. However, if the electron is bound to an atomic nucleus it cannot have anywavelength, thus it must form a standing wave around the nucleus (see figure 5.3).This gives rise to certain allowed orbitals and in between them are other forbiddenzones. The allowed orbitals have discrete electron wavelengths, so the electrons canonly have certain discrete values of energy (Cohen-Tannoudji et al., 1977).

If a photon, which carries with it some energy, comes close to an atom, there are twopossible scenarios. If the photon’s energy matches the energy difference between twoallowed orbitals in the atom, it can be absorbed by the atom and excites the electronto another orbital. If the photon’s energy does not match any energy difference, thephoton is not absorbed (Cohen-Tannoudji et al., 1977).

This explains the absorption spectrum of elements, that only the photons with acorrect energy will be absorbed and they correspond to the missing colors in thespectrum. Different elements have different orbitals, so each element will have itsunique set of discrete orbital and thus a unique absorption spectrum (Haken & Wolf,2005).

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5 Atmospheres outside our solar system

Figure 5.3: An electron wave around an atomic nucleus. On the left, the electronforms a standing wave and is thus in an allowed orbital. On the right it does not,so it is in a forbidden orbital.

Atomic orbitals are not the only place that discrete energy levels can be found inquantum mechanics. It is even less intuitive, requiring more sophisticated math-ematics, yet it can be shown that molecules can only rotate with some discreterotational velocities. Furthermore, chemical bonds in molecules can be stretch backand forth, but only with some discrete vibrational frequencies. Different bonds vi-brate with different frequencies. The bond stretching spectral lines are not as sharpas the electron excitation lines, since the vibrational frequency depends on manyparameters, such as the electron configuration of the bonded atoms and if there areother atoms nearby, which distort the bond. Thus, infrared spectra of moleculeshave wider and “noisier” spectral lines, but it is still possible to determine whichbonds are present in the material under study (Ágúst Kvaran, 2013). One suchspectrum can bee seen in figure 5.4.

5.3.2 Complications

Spectroscopy is useful in determining the composition of the atmosphere of an exo-planet. There are two main difficulties to overcome: A good light source to measurethe spectrum is needed and determining which part of the signal is really due tocompounds in the planetary atmosphere.

When using spectroscopy in a laboratory, the experimentalist can easily choose whatkind of light source he uses to light up his sample. We do not have this luxury whenmeasuring an exoplanet light years away from Earth. The most obvious light sourceto use is then the star which the planet orbits. For that to be possible, the planetmust align itself in such a fashion that the planet is between the observer and the

40

5.3 Detection of exoplanetary atmospheres

Figure 5.4: An infrared spectrum of a chemical compound. Each bond type hasits characteristic region of absorption (Hunt & Spinney, 2009).

star. Thanks to the Kepler space observatory, we know of many solar systems alignedin this manner. As previously stated, Kepler uses the transit method to discoverexoplanets.

The usage of the planet’s parent star as a light source is a point in favour of infraredspectroscopy surfaces. The most common stars in the Universe are Red dwarf stars,long-lived stars with low mass and surface temperature. These stars emit most oftheir light as infrared light, with much less visible and ultraviolet light. It is thuspossible to use infrared spectroscopy for planets orbiting Red dwarfs, but it wouldbe more difficult to use other spectroscopy types (Carrol & Ostlie, 2007).

When a good candidate to measure a transiting planet around a decently brightparent star has been found, there can be problems regarding which parts of thespectrum result from the planetary atmosphere. The parent star’s light will notbe evenly distributed between all wavelengths. It will have some brighter emissionlines and also dark absorption lines due to absorption form elements and chemicalcompounds in the colder upper layers of the stars. Dust between us and the exo-planet under study can also absorb some of the light. The biggest problem though,is our own atmosphere. It contains all of the compounds we are looking for in exo-planetary atmospheres and they will add their “fingerprint” to the spectrum, givingfalse positive signals. To overcome this, the telescopes used must either be placedon high mountain tops where the atmosphere is relatively thin or be launched intospace, like the Spitzer space telescope. Spitzer only measures infrared light and has

41

5 Atmospheres outside our solar system

successfully been used to identify exoplanets (Ágúst Kvaran, 2013).

The spectra from exoplanetary atmospheres will then be far from ideal, as there willbe significant noise due to non-ideal lighting and interstellar dust. Sophisticatedstatistical methods must be used to extract a meaningful signal from the noise.

5.3.3 Results from observations

The earliest analysis of an exoplanetary atmosphere was conducted fairly recently.In 2001, astronomers analyzed spectral data of the exoplanet HD 209458b and char-acterized absorption lines in its visible light spectrum. This first analysis only foundevidence of sodium in the planet’s atmosphere, yet it was nonetheless a big leap inthe classification of exoplanetary atmospheres. HD 209458b marks more milestonesin the discovery of exoplanets, because it was the first planet detected using thetransit method (Charbonneau et al., 2002).

Later analysis of the IR spectrum HD 209458b revealed that its atmosphere addition-ally contains oxygen, carbon compounds and water vapour. The planet is, however,a hot Jupiter planet and therefore it is unlikely to be habitable (Charbonneau etal., 2002).

In 2008, the Hubble space telescope detected for the first time an organic moleculein an exoplanetary atmosphere. Methane was detected in the atmosphere of theplanet HD 189733b, another hot Jupiter planet (Swain et al., 2008).

As of 2011, around 50 exoplanetary atmospheres had been studied, yet none be-longed to an Earth-like planet. This is mainly due to the fact that the transitmethod is biased towards finding larger planets. With better equipment, however,the same spectroscopic method can be applied to smaller, Earth-like planets (Char-bonneau, 2012).

With the launch of the James Webb Space Telescope in 2018, astronomers hopeto be able to detect atmospheres of Earth-sized exoplanets. Budget issues havediminished these possibilities, as the Terrestrial planet finder mission of the JWSTwas sadly cancelled in 2011 (Loeb & Maoz, 2013).

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6 Terraforming Mars

The human race is curious in its nature and has just started its exploration of theUniverse and its travel away from Earth. The main focus of the scientific communityin astrobiology today is to look for places where life could and might have beganand evolved. In the future, it may be desirable to transform planets and moons to ahabitable state that could support life and humans, a process called terraformation.By those means, humans could colonize and expand to other planets and moons. Itis unknown where the best future habitat is but with modern technology, Mars maybe the best choice.

There are many obstacles for making Mars habitable for humans, but there areseveral proposed methods technologically possible. It would take a significant con-tribution of time and finances.

6.1 Main problems

6.1.1 Low gravity

Mars does not have the same gravity as Earth so humans today could not surviveon Mars. There would be physical effects of a prolonged low-gravity environment,including eyesight loss and bone and muscle mass loss. The atmospheric pressure isso low that if Mars would get warm enough to melt the icecaps the water could notexist in liquid form. The water turns directly into gas through sublimation (GravityHurts , 2001).

6.1.2 Lack of magnetosphere

The magnetic field of Earth protects it and mitigates solar radiation. Mars lacksa strong magnetosphere which is thought to be the reason for its thin atmosphere.Solar wind interacts with the atmosphere on Mars, ejecting atoms into space. Solar

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6 Terraforming Mars

wind might also rip the atmosphere away from the planet as becomes trapped inbubbles of magnetic fields called plasmoids (Solar wind ripping chunks off Mars ,2008).

Comparing this to Earth, water stays here on Earth because the ionosphere is perme-ated with a magnetosphere. Hydrogen ions present in the Earth’s ionosphere movevery fast due to their small mass, but are unable to escape to outer space becausetheir trajectories are deflected by the Earth’s magnetic field. Earth’s ozone layerprotects the H2O atoms from separating to ions, due to UV radiation Terraformingof Mars (2013).

Early Mars was once believed to have a magnetic field and perhaps an atmospherethat could hold liquid water. Magnetic analysis of the Martian surface indicatesthat the magnetic field withered away for about 4.5 billion years ago. Withoutthis shield, streams of ionizing particles from the Sun would strip away a planet’satmosphere. This would have killed any life that may have emerged or forced itunderground Zala (2009).

One theory is that the Martian magnetic field disappeared during the Late HeavyBombardment. This theory fits with the observation, as only the oldest impactcraters on Mars are magnetized. Mars was hit by at least five particularly largeasteroids during the bombardment. Any one of the super-giant impacts could haveshut off the Martian magnetic field due to the heat release. Earth likely sufferedthe same bombardment impact like Mars did, but having almost twice the radius ofMars, it withstood and recovered from the huge impacts Zala (2009).

Many studies shows that Mars might also recover with time. The impact that hitMars 4.5 billion years ago may have crippled it’s magnetic field, but not destroyedit outright. In the future, it might regenerate and thus a stronger magnetic fieldcould form on Mars in the future (Grossman, 2011).

6.2 Advantages

There may be problems with terraforming Mars but on the other hand there arealso several advantages. According to modern theories, Mars in on the border of aregion known as the extended habitable zone, where greenhouse gases could supportliquid water on the surface at a sufficient atmospheric pressure (Kaufmann, 2013).

The Martian environment was once believed to be similar to Earth in its early stages.Curiosity confirms evidence of habitable ancient Mars and water appears to haveonce existed on the surface, possibly only at the poles (Kaufmann, 2013).

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6.3 Terraformation

The atmosphere of Mars contains many important elements needed for life. Mostof the oxygen in the Martian atmosphere is in the form of carbon dioxide, the mainelement of the atmosphere, as seen on figure 6.1.

Figure 6.1: The percentages of the five main gases in the atmosphere of Mars,measured by the SAM instrument (The Five Most Abundant Gases in the MartianAtmosphere, 2012).

6.3 Terraformation

Terraforming Mars would require several main processes: building the atmosphere,warming the planet substantially to 290 K and increasing the atmospheric pressurefrom 6 mbar to 61.8 mbar, which is required for people to live in. An atmosphericpressure of 10 mbar is required for plants to live in. The atmosphere must also beprevented from being lost to outer space and volatiles necessary for life must beadded to it.

6.3.1 Increase temperature

The atmosphere is mainly made of CO2, a known greenhouse gas. When the planetstarts heating up, CO2 may help with keeping the thermal energy near the surface.When this occurs, more CO2 should enter the atmosphere from the frozen ice capson the poles, strengthening the greenhouse effect. This means that building the

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6 Terraforming Mars

atmosphere and heating the planet would affect each other and Mars would beginto terraform Terraforming of Mars (2013).

Large amounts of water are stored at Mars’s south pole. If melted, it would createan ocean of up to 11 meters deep (Zubrin & McKay, 1995).

Figure 6.2: Dynamics of the Martian polar atmosphere. Current equilibrium isat point A.Raising the temperatures of the poles by 4K would drive equilibrium Aand B together causing runaway heating that would lead to elimination of the southpolar cap (Zubrin & McKay, 1995).

6.3.2 Introduction of life

After having increased the temperature and gaining little water on Mars, the nextstep would be introducing life to Mars.

In April 2012, it was reported from the Mars simulation laboratory, maintained bythe German Aerospace center, that lichen and cyanobacteria survived and showedremarkable adaptation to the same atmospheric, temperature, radiation and pres-sure conditions that they would experience on the Martian surface (Baldwin, 2012).

At this point, microbiologic life could maybe be introduced on Mars. Increasingoxygen levels in the atmosphere would require having to simulate the origin of oxygen

46

6.4 Comet impact

on early Earth (Biello, 2009). Cyanobacteria and algae could increase oxygen levelsin the Martian atmosphere by photosynthesis (Hsu, 2010).

6.3.3 Introducing nitrogen

Increasing nitrogen levels on Mars may be the biggest problem of terraformation. Asseen on figure 6.1, there is almost no nitrogen in the Martian atmosphere. Nitrogenlevels in the Earth’s atmosphere are very high. Living organisms require nitrogento form key amino acids. It is believed that the origin of nitrogen on Earth camefrom meteoroids and volcanic outgassing, as mentioned in section 3.1.

It is possible that large amounts of nitrogen exist in frozen form on asteroidal objectsin our solar system and it may be possible to move such objects in the right directionwith nuclear bombs Mars - Terraforming Wiki (2013).

Many ideas float around that are science fiction, but there remains the issue ofgravitation. Mars’s gravity is sufficient to hold onto an atmosphere for thousands,even millions of years. Many people question if this is long enough for human life.Once terraformation is complete, a habitable Mars may exist for many generationsto come (Kaku, 2010).

6.4 Comet impact

In January 2013, a new comet was discovered by astronomers in South Wales, Aus-tralia with the Uppsala Schmidt Telescope at Siding Spring Observatory. Whilehighly unlikely, there is a chance that this comet could collide with Mars on the19th of October, 2014. The trajectory of the comet 2013 A1 is still not very wellknown.

On the 19th of October 2014, the comet is forecast to possibly reach an apparentmagnitude of -8 to -8.5, as seen from Mars, making the comet 15 to 25 times brighterthan Venus. The nucleus is estimated to be anywhere from 8 km to 50 km indiameter.

If collision occurs, it can be estimated that the impact resulting from a velocity of56 km/sec could create an impact crater on Mars up to ten times the diameter ofthe comet’s nucleus and up to a 2 km deep “hole”. The energy released is equivalentto 2 · 1010 megatons (Rao, 2013).

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6 Terraforming Mars

Rough calculations show the collision to be 25 million times larger than the largestnuclear weapon ever tested (Plait, 2013). This large comet will have devastating

Figure 6.3: Diagram showing the orbit of Mars and the estimated orbit of C/2013A1 (Brown, 2013).

effects on Mars, since a collision with C/2013 A1 would result in a radical transfor-mation of Mars. The impact would be great enough to release frozen carbon dioxide,changing the atmosphere on Mars. It could either enhance the greenhouse effect orcause a decrease in surface temperature, due to dust storms preventing sunlight fromheating the planet. Volcanic activity may possibly be triggered at the impact site.

If the comet passes Mars untouched, it may still affect the atmosphere on the redplanet due to the gaseous coma, 100.000 km in diameter, that Mars would passthrough (Novosti, 2013).

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7 Conclusion

The formation and evolution of stars, planets and moons and their atmospheresare not a random occurrence, rather they follow physical laws. Science has beenconsistently searching for answers both within and outside our solar system. Life,as we know it, is possible on Earth due to dynamic processes occurring progressivelyover the past 4.6 billion years. The development of planetary atmospheres includesa delicate balance of chemical, biological and geological occurrences. Events suchas comet impacts, thermal and non-thermal escape and impact erosion interrelatesimultaneously with a planetary bodies’ composition, mass, escape velocity andsurface temperature. Outgassing, tectonic activity and whether a magnetic fieldexists, determine the content and the depth of an atmosphere and the processesthat take place in it. Exoplanetary detection methods are opening up new depths ofworlds and understanding beyond not so recent borders. The search of life on otherplanets will help us to better understand where we came from. This knowledgemight also open up the possibility of making other planets, or moons, habitablein the future. Unfolding the mysteries of life and the planetary atmospheres thatsupport our species and possibly, other species on exoplanets, is a complex yetinteresting journey.

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