CHAPTER 1: Origin of the Universe and Planetary Systems

According to our current understanding, the Universe starts roughly 13.7 billion years ago. The earliest phases of the Big Bang are subject to much speculation since astronomical data about them are not available. In the most common models, the Universe was filled homogeneously and isotropically with a very high energy density and huge temperatures and pressures, and was very rapidly expanding and cooling.

In the first 10⁻⁴³ seconds of its existence, called the Planck epoch, the Universe was very compact. It's thought that at such a dense energetic state, the four fundamental forces — gravity, electromagnetism and the strong and weak nuclear forces — were forged into a single force. The Planck epoch was succeeded by the grand unification epoch, during which gravitation separated from the other forces as the Universe's temperature fell. At about 10⁻³⁷ seconds, a phase transition occurred, during which the Universe grew exponentially and temperatures dropped by orders of magnitude. Microscopic quantum fluctuations that originated from the Heisenberg uncertainty principle were then amplified into the seeds that would later form the large-scale structure of the Universe. The model of breakneck expansion, called inflation, may explain why the Universe has such an even temperature and matter distribution.

At a time around 10⁻³⁶ seconds, the electroweak epoch begins when the strong nuclear force separates from the other forces, with only the electromagnetic force and weak nuclear force remaining unified. It is possible that some part of this decay process violated the conservation of baryon number which resulted in the predominance of matter over antimatter in the present Universe.

Another phase transition put the fundamental parameters of elementary particles and interactions into their present form, with the electromagnetic force and weak nuclear force separating at about 10⁻¹² seconds. Afterwards, the picture becomes less speculative since particle energies drop to values that can be attained in accelerator experiments.

As temperature was dropping, the quarks and gluons could combine to form baryons and antibaryons. These mutually annihilated when the temperature was even lower leaving a small fraction of original protons and neutrons due to the small excess of matter over antimatter. A similar process happened at about 1 second for electrons and positrons. After these annihilations the energy density of the Universe became dominated by photons.

The Universe continued to expand but at a much slower rate. As time passed and matter cooled, more diverse kinds of particles began to form, and they eventually condensed into the stars and galaxies of our present Universe.

The elements observed in the Universe were created in either of two ways. Light elements (deuterium, helium, and lithium) were produced in the first few minutes after the Big Bang during the primordial nucleosynthesis (or Big Bang Nucleosynthesis, BBN). Heavier nuclei are thought to have been produced much later in the interiors of stars or by cosmic rays.

The standard picture of primordial nucleosynthesis assumes that the Universe expanded according to the laws of general relativity at a given homogeneous baryon-to-photon ratio and contained only standard model particles. In the nutshell, it can be described as follows.

In the very early Universe, the temperature was so high that all matter was fully ionized and dissociated. As the Universe expanded and cooled, the interaction rates between particles declined and different particles departed from thermal equilibrium at different epochs, depending on their interactions.

At high enough temperatures (typically a few MeV), the Universe was composed mostly of photons, light leptons, and some free baryons. The charged-current weak interactions interconverted neutrons and protons maintaining them close to their thermal equilibrium value. At t ∼ 1 second after Big Bang, the temperature dropped to T ∼ 1 MeV, slightly less than the neutron-proton mass difference, and the weak interactions between neutrons and protons became too slow compared to the expansion rate to maintain equilibrium. Neutron and proton abundances froze-out at neutron-to-proton ratio n/p = 1/6. Lighter protons were more abundant as they were favored over more massive neutrons. This ratio continued to decline slowly due to the neutron β−decay (n → p + e + ν̄_e) and dropped to 1/7 by the time first nuclear reactions began.

Owing to the low density of the primordial plasma, the production of light elements occurred primarily by neutron or proton capture while 3 and 4-body reactions were suppressed. Thus, the synthesis of ⁴He and other elements could not start before deuterium was created with sufficient amount. This occurred once the temperature dropped below ∼100 keV (t ∼ 2 min after Big Bang) and deuterium could be synthetized in the process n + p → D + γ without being immediately destroyed by the energetic background photons. The abundance of deuterium in the Universe increased rapidly at that moment. As the neutrons got exhausted, the deuterium production stopped and its abundance started to decrease as it was burned to ³He and ⁴He. However, some fraction of deuterium was not converted to heavier nuclei resulting in the relic deuterium to hydrogen fraction of order of 10⁻⁵.

The nuclear reactions quickly incorporated all available neutrons into ⁴He at about 200 seconds after Big Bang with the final mass fraction of about 0.25, a number limited by the availability of neutrons at BBN. The burning of ³He into ⁴He during BBN was incomplete due to a sharp decline of ⁴He fusion rate with decreasing temperature, leaving the ³He abundance at the end of BBN of order of 10⁻⁵, similarly to deuterium. ³H which was formed along with deuterium decayed into ³He with a 12-year half-life, so no ³H has survived to the present. Because ⁴He is the most tightly bound of the light nuclides and because there are no stable mass-5 nuclides, to jump the gap to higher masses requires the Coulomb-suppressed reactions of ⁴He with D, ³H, or ³He. This guarantees that the abundances of the heavier elements are orders of magnitude lower than that of ⁴He.

A small amount of ⁷Li was produced via radiative capture ⁴He + ³He → ⁷Be + γ with ⁷Be as an intermediate state which decayed into ⁷Li via electron capture with a half-life of 53 days. In the early Universe the electron capture could not effectively proceed until the temperature dropped sufficiently low, shortly before the hydrogen recombination, with a total yield estimated to about ⁷Li/H ∼ 10⁻¹⁰. Finally, there is another gap at mass-8, ensuring that during BBN no heavier nuclides were produced with astrophysically interesting abundances. An interesting analytic approach of the primordial nucleosynthesis is presented in.^1,2 

The outcome of the standard BBN depends on a set of parameters: neutron mean lifetime which sets the relative number of neutron and protons at the onset of BBN, number density of baryons governing the reaction rates during BBN, and the number of neutrino flavors which influences the expansion rate of the Universe. The precision of measurements of the neutron mean-life has greatly improved during recent years.^3,4  The baryon density can be determined with a great accuracy by Cosmic Microwave Background (CMB) experiments⁵  and the number of neutrinos is determined from the Z boson width in accelerator experiments.⁶  Also, the nuclear reaction rates have received particular attention.³  These allow one to make firm predictions for the light element abundances and compare them with observations.

The deuterium abundance is an excellent indicator of the baryon density as BBN is the only significant known source of deuterium which is then burned into helium in the stars owing to its very weak binding. Thus, traces of deuterium observed anywhere in the Universe in the post-BBN epochs provide a lower bound to the primordial deuterium abundance. High resolution absorption lines in the spectra of distant quasars by the low-metallicity clouds of intergalactic gas reveal the presence of deuterium in agreement with BBN predictions.^3,7,8 

The experimental verification of the ⁴He abundance is complicated by the fact that it can be synthetized in stars. Nevertheless, the measurements of ⁴He abundance in ionized metal poor regions (HII) provide a good estimate of about 0.24672 ± 0.00017 in mass fraction^5,9 –in quite good agreement with the BBN predictions.^3,7,10  3He can be observed in emission lines from regions of ionized gas. The ³He nucleus has a net spin, so the analogue of the 21 cm spin-flip transition in neutral hydrogen occurs at 3.46 cm, providing the observational signature of singly ionized ³He. However, contrary to ⁴He, ³He can be both produced and destroyed in stars, thus its abundance is not well constrained by those observations.⁶ 

⁷Li is the most problematic among all the light elements.^7,11  The lower bound on ⁷Li abundance is deduced from old metal poor stars in the Galactic halo where the lithium abundance is almost independent of metallicity, forming the so-called Spite plateau. This plateau is interpreted as the BBN yield, assuming that lithium has not been depleted at the surface of the stars. There is a tension, or rather a real conflict, between the observational value for ⁷Li abundance and the theoretical prediction of standard BBN calculations which gives a value about 3 times higher. The proposed solutions to the “lithium problem” include mechanisms for reinforced depletion (by mixing or diffusion) of ⁷Li in the atmospheres of stars,^12,13  enhanced ⁷Be destruction,¹⁴  some non-standard BBN scenarios such as inhomogeneities of neutron-to-proton and baryon fraction at small scales allowing for a large variation of the element abundances,¹⁵  a beyond standard model physics such as decays of supersymmetric particles^16–19  or time variations in the fundamental constants.²⁰ 

The final abundances of the BBN elements are fundamental for the dynamics of subsequent processes of recombination of neutral atoms and formation of primordial molecules during the Dark Ages.

After BBN the Universe was still too hot for the formation of neutral atoms. The radiation and matter (electrons and light nuclei) remained in thermal equilibrium. At redshift z ≈ 40 000, when the temperature became low enough, the recombination of electrons and nuclei started by Li³⁺ + e⁻ → Li²⁺ + γ followed by Li²⁺ + e⁻ → Li⁺ + γ. The lithium ions Li⁺ did not recombine entirely due to the low ionization energy. Helium recombined first at redshift z ≈ 6000 via reaction He²⁺ + e⁻ → He⁺ + γ followed by He⁺ + e⁻ → He + γ at z ≈ 2500. Unlike lithium, helium recombined almost completely because many free electrons were available, as long as hydrogen atoms were mostly ionized. Finally, hydrogen recombined at z ≈ 1300 via H⁺ + e⁻ → H + γ.²¹ 

During recombination, the number density of free electrons fell sharply (but not instantaneously) and at z ≈ 1100 radiation decoupled from matter i.e. photon mean free path, determined by the Thomson scattering by free electrons, became longer than the horizon distance. The Universe became transparent to radiation which we can observe today as the cosmic microwave background of temperature 2.725 K⁵.

Just after the recombination, the Universe was filled with neutral atoms, some small fraction of ionized nuclei and free electrons and a large number of free photons. Photons black body spectrum peaked in infrared at that time but still the contribution of visible photons was important. As the Universe expanded, the black body spectrum was entirely shifted to the infrared. There were no stars yet and thus the Universe appeared to be completely dark. During this dark epoch, which lasted for millions of years, many chemical reactions took place forming the first molecules.

The first molecules and molecular ions must have been formed in the absence of any surfaces of dust grains, which were created after the first stars appeared. Thus, the chemistry of the early Universe is far different from the typical astrochemistry of the interstellar medium. In the low-density gas after recombination, 3-body reactions were inefficient, leaving only radiative processes to form the first molecules, mostly through radiative association. The main reactions governing the abundances of each species are described in the literature and numerical studies include more than hundred reactions and tens of atomic and molecular species in the computation of the abundance of the first molecules.^21–26 

Helium is the first neutral atom to recombine and the first atom to form a molecular bond in the early Universe. The first reactions to occur, when the hydrogen was still ionized and the helium partially ionized, were H⁺ + He → HeH⁺ + γ and He⁺ + He → He₂⁺ + γ. As the residual ionization of helium is very low, He₂⁺ was produced in substantial abundances only during the epoch of helium recombination reaching the maximal abundance of about 10⁻²², but its production became negligible after that. Moreover, the energy of CMB photons was high enough to destroy the He₂⁺ by photodissociation. The HeH⁺ could have been produced during all the epochs until z ≈ 20 because of the incomplete hydrogen recombination. The main reactions included photodissociation by CMB photons at z > 300, collisions with hydrogen atoms at 300 > z > 20 and dissociative recombination at z > 20.

The first neutral molecule to form was molecular hydrogen H₂. Radiative associations of H and H⁺ (H⁺ + H → H₂⁺ + γ) produced the molecule via H₂⁺ + H → H₂ + H⁺. The abundance of H₂ rose rapidly until z ≈ 800 when the ionization fraction of hydrogen dropped to 1% and the radiative associations with H⁺ become much less efficient. Various reactions described in detail in literature contribute to the production of H₂ during later epochs until the freeze-out at z ≈ 10 with the fractional abundance of about 10⁻⁷ relative to the total number of baryons. This way, molecular hydrogen became the most abundant molecule in the early Universe.

Deuterium is also important in the early Universe mainly because HD molecules have non-zero dipole moments and can potentially cool the primordial gas through dipole radiation to substantially lower temperatures, provided a significant abundance is obtained. The evolution of HD follows closely that of H₂ because, in general, reactions that proceed for hydrogen will proceed for deuterium as well. The freeze-out value of HD/H₂ ratio is ≈ 7*10⁻⁴ with an enhancement factor of ∼25 over the initial D/H abundance. This fractionation occurs mainly due to some additional channels of production of HD that were not allowed for pure H (for example, a direct radiative association D + H → DH + γ) and because deuterated molecules are more tightly bound than their pure hydrogen isotopes.

Another efficient coolant of the gas during gravitational collapse of first objects was LiH. After recombination of Li³⁺ and Li²⁺, the remaining Li⁺ and Li react with electrons, H, D and He atoms to form Li⁻, LiH, LiH⁺, LiD, LiD⁺ and LiHe⁺. The formation of these molecules in significant amounts is delayed to low redshifts z < 200 as they are easily photodissociated at higher temperatures. Moreover, the abundance evolution is very sensitive to the effects of nonthermal photons from recombination of hydrogen and helium and to collisions with hydrogen atoms. The chemical network for lithium is rather complicated, but the final abundances are very small ∼10⁻²⁰–10⁻²³.

The observations of most distant quasars and most ancient stars in the Galaxy point to at least some mixture of heavier elements based on C, N, O or F nuclei.²⁷  While they can be produced in the first population of stars and galaxies, their dispersion throughout the intergalactic medium points to production in the earlier epochs.²⁷  These heavier molecules could be important, if produced in sufficient amounts, for the formation of the first structures in the Universe. Even if their abundances are proven to be quite low, they could still furnish an interesting cooling agent in the gravitational collapse of proto-structures, due to their typically higher dipole.

Several authors considered the chemistry of heavier elements in the Dark Ages.^27,28  Carbon is the most likely heavy element to exist at early times. In some non-standard BBN scenarios, the carbon-based molecule CH could be produced in amounts comparable to LiH, and more abundant than HD⁺ and H₂D⁺, while the OH would be only one order of magnitude less abundant than HD⁺ and H₂D⁺.²⁸  If carbon did exist before the first stars were formed, the formation of heavier elements in the first generation of massive stars, rather than in a second generation, might also be facilitated, which could explain why all quasars and galaxies, even those seen at the highest redshifts, already exhibit spectra of heavy elements, as do the oldest known Galactic stars.

At recombination the Universe was highly uniform, whereas today it is condensed into stars, galaxies, and clusters of galaxies. How and when this occurred is one of the outstanding problems in astrophysics. We know that the first objects must have collapsed before z ≈ 5–15 as quasars are seen at that time. We also know from the observed spectra of distant quasars that the Universe had reionized by this time and that these processes were controlled by atomic and molecular physics.

The collapse of first objects was controlled by molecular cooling processes. The objects must have radiated their gravitational energy in order to collapse. Once the first galaxies formed from the collapsing gas, they started to emit ionizing radiation. The initially neutral intergalactic medium started to be ionized, first in the regions surrounding the first objects, but with time, the ionized fraction of the gas in the Universe increased until eventually the regions with the ionized hydrogen (HII regions) begin to fill a large part of the Universe.²⁹  This is called the epoch of reionization.

In the absence of heavy elements, the physics of the first star-forming systems was much simpler than that of present-day molecular gas clouds. They were made almost entirely of hydrogen and helium. As they exploded as supernovae, they ejected the heavy elements produced in their interiors into the interstellar medium. This started the cosmic chemical enrichment that led to the formation of the stars and planets that we see in the Milky Way today. Simulations and theoretical arguments predict the first stars to be substantially more massive than the chemically enriched stars that formed at later epochs.^30,31 

Gas and dust are building blocks of stars. Their composition and abundances evolve with cosmic time and are also altered by local physical and chemical processes in star formation sites. In particular, interstellar molecules play a critical role in the formation, dynamics and evolution of new stars. The study of formation, destruction and excitation of molecules – astrochemistry – encompasses astronomical observations of molecular spectra, quantum chemical computations of molecular interactions and processes on dust grains, and laboratory experiments mimicking chemical processes in space.³² 

Star formation is a continuous process which uses the available reservoir of interstellar matter and re-processes it as the star evolves. Molecules form and survive in the dense and cold regions of molecular clouds, which are shielded by dust grains from the ultraviolet radiation field (see Figure 1.1). During the collapse of the cloud under self-gravitation, molecules efficiently cool the gas, decrease its pressure, and lead to the formation of dense core(s) and ‘baby stars’ i.e., protostars, in their centers. Protostars accrete matter from the surrounding core and, to decrease their angular momentum, eject a fraction of material in the form of bipolar outflows. The mass of the star increases which leads to the ignition of the deuterium burning in its center, and the remaining material concentrates due to rotation in the protoplanetary disk. With the ignition of hydrogen burning in the stellar core, the star enters the ‘main sequence’ stage characterized by steady production of energy in the thermonuclear reactions. In the meantime, processes leading to planet formation take place in the disk. Once all hydrogen is transformed into helium in the stellar core, the stars become a ‘giant’, and enters phases with significant mass loss due to stellar winds. This way, heavy elements produced in the star's center and moved by convective motions to its surface, are injected back into interstellar medium. The abundance of elements heavier than H and He increases with the age of the Universe.

Formation of stars occurs in molecular clouds which constitute the coldest and the densest part of the interstellar medium,³³  with kinetic temperatures T_kin ∼ 10 K and number densities n_H > 30 cm⁻³. Most stars form in giant molecular clouds, which show an almost uniform mean surface density of 35 M_Sun/pc² throughout the entire Galaxy³⁴  (1 pc ∼3.26 light years ∼3.2 10⁻¹⁴ km). They are complex structures which contain dense fragments – clumps – the birthplaces of stellar clusters, and dense cores, which form individual stars or close multiple systems of stars (see Figure 1.2). The clusters are groups of physically connected stars with minimum 35 members³⁵  and minimum stellar mass density of 1 M_Sun/pc³.

Observational surveys of molecular clouds show their spatial connection to spiral arms in the Milky Way, which contain the bulk of molecular gas,^36,37  but are likely disrupted in the inter-arm space.³⁸  Clumps show a more uniform distribution, with molecular gas fraction and star forming rates similar in the spiral arms and the closeby interarm regions.^39–41  For that reason, clumps are useful tracers of large-scale star formation in the Galaxy which is subject to environmental changes as a function of the Galactocentric radius, R_GC (the distance from the Galactic Centre). For example, there is a well-known decrease in metallicity of the gas with R_GC.^42,43  Low metallicity translates to the decrease in the abundance of dust and of the ratio of CO and H₂, thus lower shielding from radiation and the cooling of the gas. As the radiation field decreases with R_GC as well, and is accompanied by several other physical effects, the decrease of star formation rate is not as high as expected.⁴⁴ 

Formation of individual stars is associated with many physical phenomena that occur simultaneously: gravitational collapse of a natal, dense core; gas and dust accretion toward the planet-forming disk; jets and winds driving supersonic shock waves; outflows sweeping up surrounding material; and ultraviolet radiation generated by the accreting star heating the outflow cavity walls. In the earliest phases of star formation, the interaction between the jet, wind and the dense envelope is particularly strong and produces spectacular outflows.^45–47  The central protostar, its surrounding envelope of material, and the outflows are jointly referred as a young stellar object.

The phase of active accumulation of mass of the central protostar lasts about 0.5 million years (Class 0/I).^48,49  Once the external mass reservoir is either accreted or dissipated, the stellar growth depends on the accretion from the gas-rich protoplanetary disk for ∼2 million years (Class II, see Figure 1.3). At the final stages, the disk is gas-poor and subject to final shaping of the planetary systems.⁵⁰  The star ignites H in its core and enters the stellar evolution path.

Abundant molecules such as H₂, CO, and H₂O are sensitive tracers of physical processes and conditions in star-forming regions. H₂ is most widespread, but due to high-lying energy levels, probes primarily the hottest gas associated with shock waves. Its emission suffers from dust extinction and is therefore most relevant for characterizing gas along the jets, away from the protostar. CO is the second most abundant molecule excited by collisions with H₂, well-suited to characterize outflows (e.g., their mass, energy, and momentum)⁵¹  and physical conditions of the gas (e.g., the temperature and density).^52,53  H₂O is a unique tracer of dynamics and gas properties in the outflows, as recently proved by the ‘Water in star-forming regions with Herschel’ program.^54,55  All these molecules account for the bulk of line emission during Class 0/I phases and as such are useful probes of total mechanical energy injected by the outflows to the interstellar medium.⁵³  They have also a potential to fully reveal the influence of environment on the physics of star-forming regions. Some molecules are used as thermometers, probing the temperature structure of star-forming regions (e.g., CO, H₂CO). Ions measure the ionization due to cosmic-rays (e.g., HCO⁺), whereas absolute abundances of molecules inform about the metallicity impact on the local chemistry.

Planet formation takes place in protoplanetary disks during the Class II/III phases.⁵⁶  The planet composition is expected to be inherited, at least to some extent, from the interstellar medium and molecular clouds. However, physical and chemical processes from clouds to planets and, in particular, inside protoplanetary disks may have a significant impact on the chemistry of the available building blocks.^57,58  It is also very difficult to decipher which molecules are inherited from dense molecular clouds, and which might be produced through chemical processing in the protoplanetary disk. By tracing the trail of molecules from clouds to planetary systems, it is possible to provide constraints to scenarios of planetary systems' formation and evolution. Since measurements of the molecular content of protoplanetary disks and the composition of exoplanets remain very challenging, the chemical complexity of clouds and disks is often compared to small bodies in the Solar System. In particular, comets and carbonaceous meteorites are critical messengers of the composition of our planetary system at the time of its formation and its subsequent evolution.

An excellent way to follow the trace of molecules through the formation and evolution of planetary systems is the measurement of isotopic ratios. For example, the origin of water content on the Earth is investigated by measuring the fraction of HDO (heavy water) to H₂O. It is likely that current H₂O must have been delivered to Earth by comets or asteroids, because H₂O formed in situ evaporated early on. Deuterium is not produced in stars and is solely a product of primordial nucleosynthesis (see Section 1.1), so its abundance in astronomical bodies is highly indicative of their origin. In particular, fractionation processes are very sensitive to local physico-chemical conditions. Over the last few years, a significant diversity of D/H measurements in cometary water has been shown to exist, from values consistent with the Earth's oceans, to values two or three times higher.^59–62  This indicates that at least part of the Earth's water might have been delivered through impact by comet-type bodies.

Recent instrument advances have also made it possible to measure isotopic ratios in protoplanetary disks, providing further insights into the processing of molecules through the planetary formation process.⁶³ 

Further information about the chemical complexity of cometary ices and their role in the Solar System formation has been gathered by the ESA mission Rosetta to the Comet 67P/Churyumov-Gerasimenko.⁶¹  Measurements of complex organic molecules with the ROSINA instrument show large similarities in the abundances of CHO–, N, and S-bearing species to the ones found in low- and high-mass protostars.^64,65  Some differences indicate the impact of processing of material in protoplanetary disks, which is subject to detailed modeling. Prebiotic chemicals, including the amino acid glycine, other complex organic molecules, and phosphorus have been found by the ROSINA instrument in comet 67P,⁶¹  while numerous amino acids have been found in meteorites.⁶⁶  This confirms that complex molecules form along the star formation process, or even earlier on.

The discovery of thousands of exoplanets since the 1990's has revealed an unexpected diversity of planetary systems and shown that our own solar system is far from typical. This triggered the development of new ideas on how planets and planetary systems form and evolve. The full picture of how planetary systems emerge from the protoplanetary disk surrounding young stars is a complex problem and our understanding of each step of this process is in constant evolution.

Planetesimals are considered as the seeds of planets. They are defined as bodies large enough to be dominated by self-gravity, which typically corresponds to minimal size in the 100 m to 1 km range.⁶⁷  Planetesimals are formed in the protoplanetary disk, the dense disk of gas and dust that surrounds a star after its formation. The material available to form planetesimals, and later planets, differs in different parts of the protoplanetary disk. In the disk regions closer to the star, the temperatures are high and only materials with a high condensation temperature, such as iron or silicates, are present in solid form. Farther away, temperatures are lower, and more volatile species, such as water, condense to form ices.⁶⁸  The distance at which different species can condense varies with time as the disk cools down.⁶⁹ 

While the dust and ice particles in the protoplanetary disk settle to the mid-plane, they collide with each other and coagulate to form aggregates, which then grow and compact through collisions.^70,71  In the process of growing from those aggregates to full kilometer-size bodies, planetesimals face multiple growth barriers and our picture of how planetesimals form is in constant evolution. Planetesimals are formed relatively fast in protoplanetary disks. In the solar system, there is evidence that large planetesimals formed on a timescale of the order of a few million years only.⁷² 

Most planetesimals will keep growing to form protoplanets or be aggregated by other planetesimals or planetary embryos. However, some planetesimals are left over by the planet formation process. Those planetesimals can be observed in other planetary systems as debris disks, in which constant collisions between planetesimals produce dust and gas that can be detected from the Earth.⁷³  The solar system also has debris disks: the main asteroid belt, located between the orbits of Mars and Jupiter, and the Kuiper Belt, beyond the orbit of Neptune. The asteroids, comets and Kuiper belt objects observed today in the solar system are planetesimals left over from the solar system formation process. The composition of asteroids and comets reflect their place of formation in the protoplanetary disk: icy comets formed further away from the Sun than mostly dry asteroids. They remain incredibly useful tools to study the history of the solar system.

The term terrestrial planet refers to planets with a solid surface made primarily of silicate rocks and metals, similar to the solar system inner four planets: Mercury, Venus, Earth, and Mars. In exoplanetary systems, however, terrestrial planets are likely very different from the Earth, with masses up to 10 times the mass of the Earth or more, different composition, and with or without atmospheres.

The formation of terrestrial planets starts with the formation of planetesimals in the protoplanetary disk. Once planetesimals reach a size of the order of kilometers, their gravitational attraction begins to play a significant role, and they keep growing either by colliding with surrounding planetesimals or through the accretion of centimeter to meter size objects.^74,75  Eventually, this process leads to the formation of planet embryos (typically Moon- to Mars-sized). In the solar system, planet embryos form in about 10⁵–10⁶ years only.⁷⁶ 

In the final phase of terrestrial planet formation, the gas contained in the disk has dissipated.⁷⁷  Forming planets gradually clear their surroundings of small objects and further growth is the result of accretion of other embryos, or planetesimals. In the solar system, this phase lasted about 10⁸ years. It is dominated by interactions and collisions between planetary embryos.⁷⁸  These collisions shape the architecture, mass, and potentially composition of the planets.^79,80  The Earth's moon is thought to originate from a giant impact between the proto-Earth and a Mars-size planetary embryo during this last stage of its formation.⁸¹  Late-stage accretion involves mixing and accretion of material from different radial distances and likely has a significant role in the delivery of volatiles to the terrestrial planets.

The scenario above describes the likely formation pathway of terrestrial planets in the solar system. However, the solar system is not representative of most planetary systems and exo-terrestrial planets very different from those in our solar system have been detected, among which are super-Earths. Similar processes as those described above govern the formation of terrestrial planets around other stars but additional mechanisms such as orbital migrations and dynamical instabilities play an important role in shaping terrestrial exoplanets.

The composition of a planet and its atmosphere is strongly influenced by where in the disk it formed. In the inner part of the disk, refractory materials are the only molecules that will condense.⁶⁸  In most cases, the timescale to fully form terrestrial planets is longer than the lifetime of the nebular gas at those distances (typically a few million years⁸² ), preventing terrestrial planets from accreting nebular gases. Impacts of planetesimals (or planet embryos) formed farther away in the disk onto the forming terrestrial planets during the late stages of the planetary formation process would provide a way to deliver volatiles to those planets. In systems with giant planets, the orbital properties and time of formation of the giant planets have a significant impact on the amount and composition of volatiles that can be delivered to the terrestrial planets and its source region. Because of that, a wide variation in the water content and atmospheric composition in different planetary systems can be expected.

Given its place of formation around 1 au from the Sun, the Earth is thought to have formed from the accretion of mostly dry planetesimals and embryos.⁸³  Evidence from noble gases in the present-day atmosphere indicate that some nebular gas might have been accreted by a planetary embryo later incorporated into Earth⁸⁴ , but most of this primordial atmosphere was subsequently lost through impacts and hydrodynamic escape. The Earth as we currently know it is relatively rich in water and has a significant atmosphere, the total mass of water on Earth being estimated to about 2.8 × 10⁻⁴ Earth masses.⁸⁵  Earth's water and atmosphere were thus probably delivered through the accretion of planetesimals or embryos formed farther away from the Sun where water could condense, such as in asteroids or comets.^86–89 

Because of their rapid formation, high mass, and large angular momentum, giant planets have a critical influence on the architecture of planetary systems. Giant planets in the solar system, Jupiter and Saturn in particular, contain large amounts of hydrogen and helium, indicating that they must have formed relatively fast, while nebular gases were still present.

In most cases, the early phase of the formation of a giant planet is similar to terrestrial planets: the formation of a planetary embryo.⁹⁰  However, in the outer part of the protoplanetary disk, the temperature is cold enough that some gases (among them water) condense into ice and the planet embryos are more massive and form faster.⁹¹  Because of that, they can accrete large amounts of gas from the protoplanetary disk before it dissipates and form giant gas planets.^92,93  At larger distances from the star, due to a lower density of material, the timescale to form a planet core by accretion of planetesimals might be too long to allow significant gas accretion and formation of a gas giant. Yet, giant exoplanets have been discovered on wide orbits.^94,95  Those would form when gravitational instabilities in the disk surrounding the young star cause it to fragment and form clumps of material. These clumps then contract and collapse to directly form giant planets.⁹⁶  Both mechanisms probably coexist to produce the diversity of masses, radii, and orbits of known giant planets and exoplanets.

The solar system has 4 giant planets: Jupiter, Uranus, Saturn, and Neptune. Jupiter and Saturn are composed mainly of hydrogen and helium, the most abundant compounds in the solar nebula. Other compounds such as ammonia, methane, or sulfur and phosphorus compounds are present in trace amounts. In addition to the nebular gases, the giant planets also incorporated rocky material and heavy elements such as iron and nickel, which form the core of the planets. Uranus and Neptune formed farther away from the Sun and even though the composition of their atmosphere is similar to that of Jupiter and Saturn, they incorporated larger amounts of ices, such as water, methane and ammonia.

The first exoplanet discovered around a main-sequence star revealed the existence of an unexpected type of planet: gas giants orbiting very close to their stars, often called hot-Jupiters.⁹⁷  Other unexpected types of exoplanets are the so-called super-Earths and mini-Neptunes, which are found around about 50% of solar-type stars.⁹⁸  Super-Earths and mini-Neptunes are planets more massive than the Earth (but less than gas giants) orbiting around their stars with periods smaller than about 100 days. Current models suggest that exoplanets such as hot-Jupiters, super-Earths, and mini-Neptunes do not form in situ and that planet migration plays an important role in shaping the architecture of planetary systems.

Planetary migration is directed inward in most cases but some mechanisms can also cause planets to migrate outwards. Planet migration can have several causes: gravitational interactions between growing planets and the gas disk surrounding it,^99–101  dynamical instabilities leading to planet–planet scattering,¹⁰²  or gravitational interaction with planetesimal disks, which played an important role in the evolution of our solar system.¹⁰³ 

Evidence of the dynamical history of the solar system can be found in the current structure of the Kuiper Belt,¹⁰⁴  or in the mixture of rocky and icy objects found in the main asteroid belt.¹⁰⁵  The latest solar system evolution models indicate that it might have experienced two major episodes of planetary migrations that shaped its architecture. In the ‘Grand Tack’ scenario, within the first few million years of the solar system, Jupiter migrated inward, before entering in resonance with Saturn and migrating outwards.¹⁰⁶  About 500 million years later, gravitational interactions between the outermost giant planets and a disk of icy planetesimals caused a scattering of icy planetesimals towards the inner solar system, while Saturn, Uranus and Neptune migrated outwards. Planetesimals gravitationally interacting with Jupiter were ejected from the solar system or propelled on highly elliptical orbits while Jupiter migrated inwards.^107–109  This model is consistent with the moon cratering record suggesting that it suffered an intense bombardment (called ‘Late Heavy Bombardment’) about 600 million years after its formation.¹¹⁰  These episodes of planetary migration are thought to have played a significant role in the delivery of volatiles to the early Earth.

In today's Universe, we observe a variety of astronomical objects from large to small scales, from galaxies and galaxy clusters to stars, planetary systems, asteroids, and comets. All these objects differ in chemical composition and physical properties, but, according to the currently accepted standard cosmological model, they have evolved over billions of years from an originally homogeneous and isotropic Universe.

The chemical composition of the observable Universe is dominated by hydrogen and helium, which were formed shortly after the Big Bang, and together account for about 99% of the chemical composition of matter. Heavier elements, such as carbon, nitrogen, oxygen or iron, constitute only less than a percent of all baryonic matter. This chemical content is continuously evolving as a consequence of reactions that take place in the star-forming regions, inside the stars themselves, and during the formation of protoplanetary disks and small bodies in the planetary systems.

Tracing the abundances of molecules in molecular clouds and star-forming regions allows the study of many important physical and chemical processes leading to the formation of new stars. On the other hand, the composition of planets, asteroids and comets reflect the place of their formation in the protoplanetary disk and remain incredibly useful tools to study the history of the Solar System.

Many observations bring continuously new insights to the occurrence, origin and distribution of chemical elements and molecules in the Universe thereby allowing us to try to explain the formation of increasingly complex molecules, including organic molecules, in distant systems. Yet many more exciting questions remain to be answered in the future.

BBN

Big Bang Nucleosynthesis

CMB

Cosmic Microwave Background