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Ewine F. van Dishoeck

Leiden Observatory, Leiden University, the Netherlands; E-mail: [email protected] and Max-Planck-Institute fur¨ Extraterrestrische Physik, Garching,Germany

Abstract. This paper presents a brief overview of the physics and chemistry in regions where new and planets are formed. New telescopes at infrared and submillimeter wavelengths reveal a rich chemistry, including simple molecules like water, complex (organic) gases, ices, polycyclic aromatic hydrocarbons, and silicates. The journey of these molecules from dark clouds to the planet-forming zones of disks is described. Ultimately, they may become part of new planetary systems where they form the building blocks for life. The continued importance of the ‘golden triangle’ of observations, models and laboratory experiments to advance the field is emphasized.

1 Introduction

When Fred Kavli looked up to the sky from his hometown Eresfjord in the mountains of Norway, he experienced ‘the world at its most magnificent’ and ‘pondered the mysteries of the Universe, the planet, nature, and man.’1 Many people across the world share Fred Kavli’s fascination with the night sky and want to know our place in the Universe, in particular the possibility for life elsewhere in space. But few people ask the question what actually lies in between the stars, and how they form or die. The space between the stars is not empty, but is filled with a very dilute gas, the so-called . The colder and denser concentrations of the gas are called interstellar clouds, and this is where molecules are detected and where new generations of stars like our Sun and planets like Jupiter or Earth are formed. With densities of 104 − 107 particles per cm3, interstellar clouds are still more tenuous than a typical ultra-high vacuum laboratory experiment on Earth. Thus, besides its astrophysical significance, interstellar space provides a unique environment in which chemistry can be studied under extreme conditions. This combination makes astrochemistry such a fascinating research field, for both chemists and astronomers alike. One of the most important developments in astronomy has been the discovery of exoplanets, i.e., planets orbiting stars other than our Sun. Although they were the subject of speculation for many centuries, the first exoplanet was found only in 1995 and study of these objects is now one of the most fast growing fields in astronomy. Today, nearly 4000 exoplanets have been discovered, and surveys have shown that on average every has at least one planet orbiting around it. This begs the questions: how did these planetary systems form? Are they similar to our ? How unique is the Solar System? And which of these planets could potentially host life, like our Earth? Thanks to a new generation of telescopes, we can now finally address these questions scientifically. Astrochemistry, also known as molecular astrophysics, is ‘the study of the formation, destruction and excitation of molecules in astronomical environments and their influence on the structure, dynamics and evolution of astronomical objects’ as stated by Dalgarno (2008) (the grandfather of this field and my PhD supervisor). This definition stresses not only chemistry but also the fact that molecules are excellent diagnostics of the physical conditions and processes in the regions where they reside. The main questions in astrochemistry therefore include: how, when and where are these molecules produced and excited? How far does this chemical complexity go? How are they cycled through the various phases of stellar evolution, from birth to death? And, most far-reaching, can they indeed become part of new planetary systems and form the building blocks for life elsewhere in the Universe?

1http://kavliprize.org/about/fred-kavli

1 Fred Kavli established prizes in three areas of interest, ’from the biggest to the smallest, to the most complex’. This lecture will touch on all three aspects: starting with the vast Universe and the gigantic interstellar clouds in , we will zoom into the smallest molecules and nanometer-sized dust particles, and end with the complex question of the origin of life. In fact, it is fascinating to realize that the macroscopic structures that we see in the sky such as clouds, stars and planets are to a large degree controlled by the microscopic processes within and between atoms and molecules. The contents of this lecture have been described in detail in recent reviews by myself (van Dishoeck, 2006, 2014a,b, 2018; van Dishoeck et al., 2013, 2014) and others Tielens (2013); Caselli & Ceccarelli (2012); Bergin et al. (2013); Herbst & van Dishoeck (2009), which contain references to the hundreds of papers that have made the progress in this field possible over the past decades. Only a few selected references will be given here.

Figure 1: Hubble Space Telescope image of the Orion star-forming region.The distance to the cloud is 1500 light years, and the image covers a region on the sky of approximately the angular size of the full moon. The dark regions are dense molecular clouds. The colors represent emission from ionized gas (hydrogen Hα and other atomic lines at optical wavelengths). Credit: NASA/StScI/ESA, M. Robberto and the Orion Treasury Project Team.

2 Birthplaces of stars: interstellar clouds

Interstellar clouds are found throughout the Universe. This talk focuses on the solar neighbourhood because that is where the sensitivity and spatial resolution of our instruments is highest. The Sun is one of several hundred billion stars in the and can be found about halfway out from the galactic center to the edge. The clouds discussed here lie at a distance of less than 1500 lightyear from the Sun (1 lightyear=9.5×1017 cm). But the same processes also happen in the rest of the Milky Way and in external galaxies, even out as far as the edge of the observable Universe. Molecules like water and carbon monoxide have been seen at distances corresponding to a time when the Universe was only 500 million years old, less than 5% of its current age. A well-known example of an is the Orion nebula, which contains a stellar nursery with hundreds of stars just in the process of forming. Beautiful images such as those from the Hubble Space Telescope or European Southern Observatory (ESO) show not only colourful nebulae due to ionized gas that is emitting brightly at visible wavelengths, but also very dark regions (Fig. 1). These dark areas are dense molecular clouds which contain tiny ∼0.1 µm-sized particles of dust composed of silicates and carbonaceous compounds. They absorb and scatter visible and ultraviolet (UV) light, thereby also protecting molecules from dissociating radiation. Molecular clouds can be quite large (tens of light years across) and massive (up to 105 solar masses) but the process of star formation is quite inefficient: only a few % of the gas is eventually turned into stars (Evans et al., 2009).

2 Figure 2: Lifecycle of gas and dust in interstellar space. Some characteristic molecules at each of the star- and planet formation and stellar death stages are indicated. Image by Bill Saxton (NRAO/AUI/NSF) and molecule pictures from the Astrochymist (www.astrochymist.org; this website also contains a list of detected molecules in space).

3 Interstellar molecules

3.1 Composition and chemistry

Interstellar clouds consist mostly of gas, about 99% by mass, with 1% of solid materials. The astronomers’ view of the periodic table is fairly restricted. Clouds in the local Universe contain primarily hydrogen (90%) with about 8% helium by number. All other elements are called ‘metals’ by astronomers even if they are obviously not a metal in a chemical sense. The next most abundant elements, important for water and organic chemistry, are oxygen, carbon, and nitrogen at abundances of only about 4.9, 2.7 and 0.7 ×10−4 that of hydrogen. Traditional chemistry would predict that virtually no molecules are formed at typical densities of 104 cm−3 and temperatures of 10 K of dense clouds, with 1000 times more hydrogen than any other chemically interesting element. The detection of nearly 200 different species over the past 45 years (not counting isotopologs) demonstrates the opposite: there is a very rich chemistry in space. Detected molecules range from simple everyday molecules like CO, + H2O and NH3 to relatively complex molecules such as CH3OCH3 and C3H7CN and unusual species like ions (HCO , + N2H ) and long carbon chains (HC7N) (see Fig. 2). Because of the low densities, the chemistry is not in thermodynamic equilibrium. Only two-body processes, i.e., reactions between two species at a time, are thought to be important and the detailed kinetics of thousands of reactions need to be considered. There are two basic processes by which molecular bonds can be formed at low densities. The first one is radiative association of atoms or molecules, in which the binding energy of the new molecule is carried away through the emission of photons. In contrast, on Earth at high densities a third body often takes away the extra energy. The second process involves formation of molecules on the surfaces of grains, in which the dust particle accomodates the released energy. Here the surfaces provide a reservoir where atoms and molecules can be stored and brought together for a much longer period than possible is in the gas. They thus enable reactions that are too slow in the gas such as the hydrogenation of atomic O, C and N to form H2O,CH4 and NH3 (Tielens & Hagen, 1982), but interstellar grains are not catalytic in a chemical sense. Molecular bonds are destroyed by intense ultraviolet radiation through the process of photodissociation (van Dishoeck, 1988; Heays et al., 2017). Molecular ions can also break up by reactions with electrons, a process called dissociative recombination which is very rapid at low temperatures.

3 Once molecular bonds have been formed, they can be rearranged by chemical reactions leading to more complex species. Reactions between ions and molecules are particularly fast at low temperatures, even down to 10 K (Herbst & Klemperer, 1973). The ions are created either by photoionization of atoms like C, S and Fe by the ambient UV radiation at the edges of clouds, or by cosmic rays, i.e. highly energetic particles consisting of protons and atomic nuclei. Cosmic rays can penetrate deep inside clouds and also ionize H, H2, He, O and N, whose ionization potentials + + are above 13.6 eV, the threshold where the interstellar UV radiation cuts off. H2 reacts quickly with H2 to form H3 , a cornerstone molecule which kick-starts the chemistry.

Ion Molecule High-T EA [K] ~10 000

+ ~3 000 O H+ O ~2 000 H 2 + H3 Surface H hv H OH+ 2 s-O s-O H2 OH 2 s-O e- 3 H O+ 2 - H, H e H hv H 2 2 hv hv O, OH H2

+ T H O+ HCO H O s-H O 3 e- 2 hv 2

Gas Phase Grain Surface

Figure 3: Simplified water chemistry illustrating the different routes to water through (left) low temperature gas-phase ion- molecule chemistry; (middle) high temperature gas-phase chemistry; and (right) surface chemistry. Figure by M. Persson, based on van Dishoeck et al. (2011).

3.2 Water as an example

An illustrative example of how different routes contribute to the formation of a particular molecule under different conditions is provided by the networks leading to interstellar water (van Dishoeck et al., 2013) (Fig. 3). At low tem- peratures and densities, the gas-phase ion-molecule reaction route initiated by cosmic rays dominates and produces a low fractional abundance of water around 10−7 with respect to hydrogen. At high temperatures such as encountered in shocks, reaction barriers of O and OH with H2 can be overcome and H2O is rapidly formed by neutral-neutral reac- tions. Finally, in cold dense clouds, formation of water ice is very efficient and locks up the bulk of the oxygen. Water can be brought from the ice into the gas phase by thermal desorption at high temperatures and by photodesorption triggered by UV radiation in cold clouds.

4 Telescopes: new windows on the Universe

Progress in astronomy is very much driven by large telescopes equipped with highly sensitive detectors. To penetrate inside dense clouds long wavelength data are needed where the scattering and absorption by dust is much less. Also, the energy levels of a molecule are quantized into electronic, vibration and rotation states, which allows unique identification of a molecule, whether on Earth or in space. Their primary ‘fingerprints’ occur at infrared and millimeter wavelengths. In particular, solids (silicates, ices) are uniquely observed in the infrared whereas the bulk of the gaseous molecules are best probed at millimeter wavelengths.

4 Figure 4: Left: Artist impression of the Infrared Space Observatory (ISO/ESA); Middle: Artist impression of the Herschel Space Observatory superposed on an Herschel image of the thermal dust emission of the Rosette molecular cloud at far-infrared wavelengths (ESA/PACS/SPIRE consortia); Right: Image of ALMA at night (ESO/NRAO/NAOJ).

4.1 Infrared space missions

In the last two decades, there have been three very important space missions at mid- and far-infrared wavelengths (2– 200 µm) (Fig. 4). This wavelength range needs to be observed from space because molecules in the Earth’s atmosphere such as water block most of the infrared radiation. The Infrared Space Observatory, launched by the European Space Agency (ESA), flew from 1995–1997, had a 60cm telescope and was equipped with four instruments. One of them, the Short Wavelength Spectrometer, was built under leadership of SRON, the Space Research Institute Netherlands (de Graauw et al., 1996). It opened up spectroscopy over the full infrared wavelength range, allowing a complete inventory of interstellar ices (see van Dishoeck, 2004, for review). Also, it demonstrated that forming massive stars like those in Orion have large amounts of warm water and carbon dioxide vapor in their immediate surroundings. Another big surprise was the detection of crystalline silicate material in disks around young and old stars, in contrast with the largely amorphous interstellar silicates. The Spitzer Space Telescope, launched by NASA, was operative between 2003 and 2009. With a cooled 85cm telescope and a new generation of detectors, it was much more sensitive than ISO and could observe also low mass stars similar to our young Sun. The ‘Cores to Disks’ (c2d) key program (Evans et al., 2003) made an inventory of ices, silicates, PAHs and gases of nearly a hundred protostars and disks in the nearest star-forming clouds. It also discovered simple molecules in the inner regions of disks (Lahuis et al., 2006; Carr & Najita, 2008). These space missions culminated with the Herschel Space Observatory, launched by ESA, which flew from 2009- 2013 and was actually the largest astronomical telescope in space at 3.5m diameter. Of particular importance was the HIFI instrument built by a large consortium of European and US institutes under leadership of SRON and consist- ing of a heterodyne spectrometer covering the 150-600 µm range where water vapor has its primary transitions (de Graauw et al., 2010). The ‘Water in Star-forming Regions with Herschel’ (WISH) program was one of the guaranteed programs, providing key insights into water at all stages of the star-formation process (van Dishoeck et al., 2011) 2.

4.2 Ground-based telescopes

On the ground there have also been significant advances. ESO’s-Very Large Telescope (VLT) and similar telescopes by other nations provide the most powerful collection of 8–10m optical-infrared telescopes on Earth. At millimeter wavelengths single-dish telescopes such as the 15m James Clerk Maxwell Telescope, the IRAM 30m telescope and the 12m Atacama Pathfinder Experiment, together with pioneering millimeter interferometers have been workhorses for astrochemistry over the past 30 years, albeit with limited spatial resolution (van Dishoeck & Blake, 1998). The major step forward in the field is the Atacama Large Millimeter/submillimeter Array (ALMA), a worldwide collaboration between Europe, North America and East Asia, located in north Chile on the Chajnantor plateau at an altitude of 5000 meter. It consists of 54 ×12m and 12×7m antennas distributed over a 16km diameter high-altitude plane, whose signals can be combined to obtain a much sharper view of astronomical objects than a single dish can.

2www.strw.leidenuniv.nl/WISH

5 ALMA became operational in 2011 and quickly proved to be transformational for the field. Its antennas can be moved to different locations on the plateau, providing a zoom-lens like capability, forming images with a sharpness and sensitivity up to two orders of magnitude higher than any previous instrument.

Figure 5: The golden triangle in Astrochemistry, linking observations, models and laboratory experiments (including quantum chemical calculations). Examples of different types of observations, models and molecular processes that have been studied in Leiden are indicated.

4.3 The golden triangle: observations-laboratory-models

Besides observations, the availability of accurate atomic and molecular data is another prerequisite for astrochemistry. Note that the term ‘laboratory astrophysics’ implies experiments as well as quantum chemical and molecular dynamics calculations. These data then feed either directly into the observations (e.g., spectroscopy) or into models of the chemical kinetics (e.g., reaction rates) that aim to explain the abundances of observed molecules and predict the presence of not-yet detected species (Fig. 5). Theoretical quantum chemistry studies can be equally important as laboratory experiments in providing information on astrochemically relevant questions. One important example is that of collisional rate coefficients, needed to deter- mine the excitation of interstellar molecules. Another prime example is photodissociation rates of small radicals and ions. Astrochemistry in Leiden has been set up along the lines of this triangle, with expertise in all three areas covered together with my colleagues Harold Linnartz, Xander Tielens, Michiel Hogerheijde and Marc van Hemert. The laboratory focus in Leiden has been on theoretical studies of photodissocation processes (van Dishoeck & Visser, 2015; Heays et al., 2017), laboratory spectroscopy, and laboratory simulations of the chemistry taking place on and in interstellar ices, with or without the presence of UV radiation (Linnartz et al., 2015).

5 Star and planet formation

Figure 2 illustrates the cycle of material from clouds to stars and planets, and ultimately back to the interstellar medium. Diffuse clouds are low-density clouds (∼ 100 cm−3) in which UV radiation can penetrate and destroy molecules. Because of the short timescales, they form the best test-beds for interstellar chemistry. Most of my early work was focused on diffuse and translucent clouds, especially on the transition of hydrogen and carbon from atomic + to molecular form (H→H2,C →C→CO) (e.g., van Dishoeck & Black, 1988, 1989). Interstellar clouds can be stable for millions of years, but eventually gravity takes over and the densest part of the cloud

6 Figure 6: Detection of ice absorption toward the low-mass protostar HH46 IRS using Spitzer data at 5–20 µm and VLT-ISAAC data at 2–5 µm (Boogert et al., 2004, 2008). The strong solid CO2 stretching band at 4.3 µm is missing since it cannot be observed from the ground. The insert shows a blow-up of the strong solid CO band; the weaker feature at 2167 cm−1 is due to OCN−. Background: Spitzer composite 3 (blue, stars), 4.5 (green, shocked H2), and 8 µm (red, PAHs) image, showing the embedded protostar with its outflow. Credit NASA/ESA/A. Noriego-Crespo. collapses to form a new star. In the standard scenario, the collapse occurs inside out so the protostar at the center of the cloud continues to grow as material from the envelope accretes onto the star. Because of angular momentum conservation, material cannot continue to fall in radially but ends up in a disk around the star where the gas is in Keplerian rotation. The young star will develop a jet and wind which can escape in a direction perpendicular to the disk. When they interact with the surrounding envelope and cloud, they create shocks and entrain material in bipolar outflows. With time (∼ 106yr), the opening angle of the wind increases and the envelope is gradually dispersed, revealing a young pre-main sequence star surrounded by a so-called protoplanetary disk. These disks are about the same size as the Solar System (∼100 AU, with 1 AU = distance Sun-Earth = 1.5×1013 cm) and contain a few Jupiter masses of 30 13 −3 gas and dust (1 MJup=1.9×10 gr or about 0.1% of the mass of the Sun). Here densities can be as high as 10 cm in the inner regions of disks. The increased densities and thus shorter collision timescales lead to gradual coagulation of dust grains to form peb- bles, rocks and planetesimals, although the precise mechanisms are not yet understood. The large particles settle to the midplane, where they can form kilometer-sized objects that interact gravitationally to form (proto)planets and even- tually a full planetary system (timescale up to 100 Myr). Comets and asteroids are remnant planetesimals, typically a few km in size, that did not end up in one of the planets and were scattered and preserved in the cold outer regions of our Solar System. Their composition thus reflects the conditions during our planetary system formation. At the end the stellar life cycle, the nuclear fuel is exhausted so the stars loose their support, swell up and start losing mass (Fig. 2). These evolved Asymptotic Giant Branch (AGB) stars are surrounded by circumstellar envelopes, i.e., dense shells of molecular material driven by radiation pressure that are a rich source of molecules. This gas, now enriched in heavy elements due to nuclear fusion, enters the interstellar medium and can become part of a new cycle of stellar birth and death.

6 Following molecules from cloud to disk

When interstellar clouds collapse to form new stars and planets, the surrounding gas and dust become part of the infalling envelopes and rotating disks, thus providing the basic material from which new solar systems are formed. Here the physical and chemical characteristics of each of these stages are discussed. Fig. 7 provides an overview of the chemical changes en route from cloud to disk.

7 Figure 7: Schematic representation of a protostellar envelope and disk with key steps in the water and organic molecule for- mation indicated. Water ice is formed in the parent cloud before collapse and stays mostly as ice until the ice sublimation temperature of ∼100 K close to the protostar is reached. Organic molecules are formed at low temperatures once CO freezes out at somewhat higher densities. Part of these ices are incorporated into the forming disk where they may become part of planetesimals and ultimately planets. Figure by R. Visser, adapted from Herbst & van Dishoeck (2009).

Early work focused mostly on high-mass star-forming regions such as Orion (Fig. 1) which have luminosities that are 104 − 105 times higher that of our Sun. Studies of solar-mass regions started 25 years ago (e.g., van Dishoeck et al., 1995), but came to fruition only in the last decade with the much higher sensitivity of Spitzer, Herschel and ALMA making it possible to study them.

6.1 Protostellar envelopes

Prior to collapse, dense interstellar clouds are cold, typically only 10 K. At densities > 104 cm−3, the timescale for < 6 collision of a molecule with a grain becomes shorter than the lifetime of the core, ∼ 10 yr, so nearly all species except H,H2 and He will freeze out and form an icy mantle on the grain. This heavy freeze-out stage is seen directly in the observations of interstellar ices (Fig. 6) and indirectly in the chemistry of remaining molecules in the gas. In particular, cold and dense gas is characterized by high abundances of deuterated molecules. The overall deu- terium abundance in the solar neighborhood, [D]/[H], is about 2 × 10−5, but the ratio of any molecule XD/XH is typically a few ×10−2 in cold clouds, i.e., enhanced by 3 orders of magnitude (e.g.,DCO+/HCO+, DCN/HCN). For molecules with multiple hydrogen atoms, even doubly- and triply-deuterated versions have been seen, such as D2CO −2 and CD3OH. The observed ND3/NH3 ratio is ≈ 10 , so the enhancement is 13 (!) orders of magnitude. This high deuterium fractionation results from two effects. First, the deuterated molecules are produced by deuteron transfer + + from H2D , whose abundance relative to H3 is enhanced at low temperatures due to the slight difference in zero-point + vibrational energy. Second, H2D is further enhanced when its main destroyer, CO, is frozen out. Indeed, the first + detection of interstellar H2D was in a cold protostellar envelope with significant CO depletion (Stark et al., 1999). Once the protostar has formed, it heats up the dust in its immediate surroundings, so the molecules locked in the ices start to return to the gas phase, roughly in a sequence according to their sublimation temperature (Fig. 7). For example, CO and N2 have much lower sublimation temperatures (around 20 K) than H2O (around 100 K), as determined by laboratory experiments. Above 100 K in the so-called ‘hot core’ region close to the protostar, most volatile organic molecules are returned to the gas phase. These evaporated species then drive a high temperature gas-phase chemistry producing second generation complex molecules for a period of ∼ 105 yr after sublimation (Charnley et al., 1992).

Hot cores are indeed observed to be rich in complex organic molecules such as CH3OH,HCOOCH3 and CH3CN,

8 Figure 8: Blow-up of part of ALMA spectrum of the deeply embedded protostar IRAS16293-2422B around 345 GHz, illustrat- ing the large number of lines detected in this source. Most lines originate from organic molecules, only a few lines from simple species are labelled. The source lies at a distance of 450 lightyears. Adapted from Jørgensen et al. (2016). both for high and low-mass protostars. ALMA now allows us to zoom into these regions on solar-system scales (∼ 30 AU radius, comparable to the orbit of Neptune in our Solar System) and make a full inventory of molecular complexity. An excellent example is the ALMA ‘Protostellar Interferometric Line Survey’ (PILS) of a young binary source, IRAS 16293-2422, in the nearby Ophiuchus molecular cloud for which more than 10000 lines have been detected (Jørgensen et al., 2016) (Fig. 8). Many complex, even prebiotic molecules have been identified in this survey, several of which had been seen previously only toward the massive Galactic Center cloud SgrB2(N). This includes the simplest sugar (HCOCH2OH, glycolaldehyde), (CH2OH)2 (ethylene glycol), molecules with peptide- like bonds (NH2CO formamide, CH3NCO methyl-isocyanate, CH3C(O)NH2 acetamide), c-C2H4O (ethylene oxide), CH3COCH3 (acetone) and tri-carbon molecules like C2H5CHO (propanal). Most importantly, these molecules are now known to be present on scales comparable with those of protoplanetary disks. They thus provide direct insight into the chemical complexity in regions where planetesimals and planets are forming. One of the surprises of the last years has been that these complex molecules are not only found in hot cores, but are seen at all stages of the star-formation cycle, even in cold clouds prior to star formation, albeit at much lower abundances (Bacmann et al., 2012). This raises two fundamental questions: how are these molecules formed without any ‘energetic’ input? The scenario sketched in Fig. 7 assumes that radicals become mobile at slightly elevated temperatures (20–40 K) to form complex molecules. And if these molecules do indeed form on grain surfaces, how are they returned to the gas at temperatures well below their sublimation temperatures? The first question has now been addressed by laboratory experiments, which demonstrate that not only CH3OH (methanol) can form on grain surfaces by hydrogenation of CO, but even molecules as complex as C3H6O3 (glyceraldehyde), as long as the relevant radicals to build these molecules (HCO, CH2OH) are located close together on the grain surface (Fedoseev et al., 2017).

Figure 9: Left: ALMA image of where CO ice (the CO ‘snow line’) has formed in the TW Hya protoplanetarydisk, as traced by + emission of the N2H molecule (Qi et al., 2013). The blue circle represents where the orbit of Neptune would be when comparing it to the size of our Solar System. It also marks the inner boundary of the region where smaller icy bodies like comets would form. Right: Detection of the ground-state lines of ortho- and para-water in the TW Hya disk with Herschel-HIFI implying the presence of at least 6000 Earth oceans of water ice (Hogerheijde et al., 2011), superposed on an artist impression of the disk (NASA/IPAC/Caltech/R. Hurt).

9 6.2 Protoplanetary disks

When the envelope has dissipated, the young star becomes visible, surrounded by a disk (Fig. 9). These circumstellar disks are the birthplaces of planets, so they are particularly interesting targets for astrochemical studies. Since disks are at least a factor of 1000 smaller than their parent clouds and contain only ∼1% of their mass, their emission is readily overwhelmed by that of any remnant cloud, however. Only with the advent of ALMA have detailed studies of disks become possible. Disks have strong radial and vertical temperature gradients due to heating by the UV radiation of their parent star, from >1000 K in the inner disk and surface layers to 10 K in the outer midplane. Thus, no single instrument or wavelength probes the entire disk reservoir: a combination of near-, mid-, far-infrared spectroscopy combined with spatially resolved ALMA data is needed. As a result, disks consist of different chemical layers (Aikawa et al., 2002): at the surface, molecules are dissociated into atoms by the strong UV radiation. Deeper in the disk, the grains are warm enough to prevent freeze-out and molecules can survive the UV radiation. Deep in the cold midplane most molecules are frozen out onto the grains. There has been enormous observational and model activity in this field in recent years (see review by Henning & Semenov, 2013). Fig. 9 illustrates two examples. Herschel-HIFI has revealed the bulk of the cold water vapor reservoir in the outer disk by very deep integrations as part of the WISH program. Since the observed water gas is produced primarily by UV photodesorption of ice, it points to the presence of a large reservoir of underlying water ice, the equivalent of about 6000 Earth oceans in the case of the TW Hya disk (Hogerheijde et al., 2011). However, many disks do not show water vapor emission down to very low limits, suggesting that by 106 yr (or even earlier) the bulk of water ice is locked up in larger bodies that do not cycle up to the warmer upper layers (Du et al., 2017). A similar situation may apply to CO in disks (Miotello et al., 2017). Another related highlight is the first imaging of ‘snow lines’ in disks, i.e., the radius where a molecule changes from being primarily in the gas phase to being frozen out as ice. Snow lines (or, more generally, ice lines) are important because they enhance the mass of solids by a factor of a few and thus facilitate planet formation just behind the snowline. Because of their lower binding energies, the ice lines of CO and CO2 are located at larger radii than that of H2O. This selective freeze-out of major ice reservoirs has the important consequence that it changes the overall elemental [C]/[O] abundance ratio in the gas and thus the composition of the atmospheres of any giant planets that are formed there (Oberg¨ et al., 2011).

The H2O snowline cannot be imaged with ALMA in protoplanetary disks since it lies at only a few AU, too small with ALMA to resolve. Moreover, cold water lines cannot be observed from the ground. However, the CO snowline at larger radii is readily imaged with ALMA. An early example for the TW Hya disk is shown in Fig. 9 (left). Rather + than using CO itself, the image shows a chemical tracer, N2H . This ion is enhanced when its main destroyer, CO, freezes out. Another tracer is the DCO+ ion.

7 From disks to comets and exoplanets

Even centuries ago, people were speculating that our Solar System was only one of many. Data from the ground and from the Kepler satellite show that there is indeed a great diversity in exoplanetary systems, none of them look like our own. What drives this diversity planet sizes and architectures? How far have planets migrated from their birth sites? Exoplanets are ultimately composed of the gas and dust from the collapsing cloud. Detecting forming planets themselves is still very challenging, but the emission of millimeter-sized particles with ALMA can indirectly infer the processes at work. In particular, once the icy dust grains in the outer part of disks grow to pebble size (few cm), they start to decouple aerodynamically from the gas and drift inward due to friction. This can happen so quickly that it is, in fact, a major question why there are any mm-sized dust grains left in disks, as seen by ALMA. A solution lies in the presence of dust traps due to gas-pressure bumps, predicted by theorists decades ago and finally found by ALMA. A spectacular example is that of Oph IRS48 (van der Marel et al., 2013), which shows a remarkably asymmetric dust structure (but symmetric gas structure) (Fig. 10) with characteristics that can only be explained by a vortex accumulating dust in a

10 Figure 10: Collection of ALMA images of protoplanetary disks showing the diversity of dust and gas structures. Typical sizes of the images are a few times that of our Solar System. Left: Gas (blue) and dust (orange) in the IRS 48 transitional disk revealing a dust vortex. Middle-left: Dust cavity in the SR24 transitional disk. Middle-right: Dust rings and gaps in the HD 169142 disk. Right: Multiple dark rings (depression in dust emission) in the young HL Tau disk. Credit: N. van der Marel/P. Pinilla/D. Fedele/ALMA Partnership-ESO/NAOJ/NRAO. trap. Inside the trap, conditions are such that dust particles can rapidly grow to planetesimal, or comet, size. Dust traps are now seen in many other disks, often in the form of dust rings or cavities, with a variety of explanations for their origin (Fig. 10). Transitional disks with large dust cavities are prime targets for ALMA since they are likely caused by planets currently forming in the disk. Altogether, it is exciting that ALMA can now reveal the location of planetesimal formation in disks –the first steps of planet formation– as well as their associated chemistry. Within our own Solar System, comets are the left-over planetesimals from its formation 4.5 billion years ago. Comets spend most of their time in the outer reaches of the Solar System where it is cold and their chemical composition can be preserved. To test whether this is indeed the case, one can compare the abundances of molecules in protostellar envelopes and disks, such as those observed with the PILS survey of the low-mass protostar IRAS16293B, with the abundances of molecules observed in comets (Mumma & Charnley, 2011). The recent ESA Rosetta mission to comet 67P/Churyumov Gerasimenko equipped with high resolution mass spectrometers (Le Roy et al., 2015) provides a new opportunity for such a comparison. Overall, many of the same molecules are detected, even in similar abundances. An interesting highlight has been the recent identification of chloromethane, CH3Cl, a molecule thought to be a possible biomarker, in both comet 67P and in the protostar (Fayolle et al., 2017). There have also been surprises, most notably the detection of abundant molecular oxygen, O2, in comet 67P (Bieler et al., 2015) but not in IRAS16293B (Taquet et al., 2018). Using the chemistry outlined in Fig. 3, this implies that our Solar System was likely formed in a somewhat warmer environment, 20–30 K rather than 10 K, to suppress the hydrogenation of most of the O2 molecules to H2O. Another intriguing question is whether comets brought the water to Earth. One option is to look at the deuterated version of water, HDO, since the HDO/H2O ratio in oceans has a characteristic value that is a factor of 8 higher than that of the Solar System as a whole: some fractionation must have taken place. Measurements of HDO/H2O by Herschel-HIFI, Rosetta and other instruments of a number of comets find ratios that range from the same value as in Earth’s oceans to values that are up to a factor of two higher. The observed values measured in protostellar envelopes, which represent ices before their incorporation into larger bodies, are in the same range. Overall, the similarity suggests that there are icy bodies in the Solar System that could have delivered water through impacts on the early Earth in its late stages of formation. Comets themselves may have contributed only a small fraction but other icy bodies from the asteroid belt may have been more effective in bringing water to Earth. Comets may, however, have been largely responsible for seeding the early Earth with organic material. Whether these relatively simple organic molecules can react to more and more complex polymers and ultimately living multicellular organisms is a key question to be answered by chemistry and biology. Life as we know it only arises in the presence of liquid water; hence, planets in the so-called habitable zone around a star are prime objects to search for signs of life in their atmospheres.

11 8 Concluding remarks

This talk has given a ‘taste’ of the many physical and chemical processes that play a role during star- and planet formation. Thanks to new telescopes, the journey of water and other molecules can now be followed from cloud to disk. Observational surveys have shown that there is a rich chemistry associated with nearly every forming star, on scales comparable with that of our own Solar System. Thus, the building blocks for life are common in interstellar space. Also, nearly all young stars are surrounded by disks with enough mass to make a planetary system. Thus, the ingredients for planet formation are also common in space. The next decades will focus on the trail from disk to planet, using new facilities such as the James Webb Space Telescope and the ground based Extremely Large Telescopes (20-40m diameter). Those telescopes will measure the molecular abundances of the inner few AU of protoplanetary disks that ALMA cannot reach, as well as the composition of the atmospheres of young and mature planets. Can we ultimately link the composition of a planet atmosphere to its formation location and thus learn about its history? The main message here is that the chemical composition of the building blocks of planets is already largely determined at the time of formation of the cloud and disk from which the star and planets originated. To continue to unravel this story of our origins, an interdisciplinary approach such as presented here will be crucial.

Acknowledgements

The research described here has only been possible thanks to the efforts by hundreds of colleagues: astronomers, chemists, physicists and instrument builders. I am particular indebted to all my PhD students and postdocs for their excellent work and creativity, and for making this topic so much fun. My research has been supported by Leiden University, the Netherlands Organization for Scientific Research (NWO), the Netherlands Research School for As- tronomy (NOVA), the Royal Netherlands Academy of Arts and Sciences (KNAW), by European Union A-ERC grant 291141 CHEMPLAN and by the Max Planck Gesellschaft.

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