Astrochemistry is the study of molecules in astronomical environments, and of their in-fluence on these environments. It encompasses astronomical observations of molecules in diﬀerent regions of space in order to understand molecular complexity, the conditions that lead to the formation of molecules, what molecules tell us about the regions they are found in, etc… but also other activities. It also includes astrochemical modelling, which takes a diﬀerent approach to these problems by trying to simulate the behaviour of the diﬀerent regions of space, based on a chemical network and a physical model of the region. The third branch of astrochemistry is part of what is called laboratory astrophysics: the experiments and theoretical calculations that bring the knowledge required to understand the observations and give inputs for the models. These three branches complement each others and none can exist without the help of the others. The field of astrochemistry, whether it concerns observations, modelling or laboratory experiments, has been extens-ively reviewed in recent years, so that it is not diﬃcult to find introductions [1, 2, 3, 4] or more detailed reviews [5, 6, 7, 8, 9, 10, 11, 12, 13] on many diﬀerent aspects. Most of this part has been written drawing from this material, although I will source in the text some specific aﬃrmations or examples.
Molecules in the universe
Molecules are found on planets, moons, comets and asteroids in planetary systems such as our own, but also in the interstellar medium, which occupies the space between stars and makes up 10% of the matter in a galaxy like ours . In the interstellar medium, molecules are found more particularly in star- and planet-forming regions. The study of molecules in space is interesting from many angles. From an « astrophysical » point of view, the emission and absorption lines of molecules in the spectrum of a given astronomical object is one of the very few information regarding that object that can be obtained from Earth. Molecules are used as tracers of some basic properties like temperature or density, but also more complex ones. But they are not just passive markers of what is happening in a given medium: they also actively participate in the physics of these regions. For example, molecules cool down dense regions by radiatively sending away energy coming from collisional excitations. From a « chemical » point of view, the study of molecules is also extremely interesting because the conditions of density and temperature and the time scales on which reactions occur are extremely diﬀerent from those on Earth, giving a radically diﬀerent perspective on chemistry. Explaining how the molecules that are observed can be formed and how far can molecular complexity go in diﬀerent regions of space is one of the main challenges of astrochemistry. One of the important motivations is also to figure out how the molecules observed in space traverse the cycle of matter, going from an interstellar cloud to a planet forming disk to comets and planets. What was the inventory of available molecules at the dawn of prebiotic chemistry, and in what way were they delivered to planets like our own, ties the study of molecules in the universe to astrobiology and the question of the origin of life.
Emphasis will be put here on the interstellar medium, in particular regions where the tiny dust grains that are part of the medium are covered with an icy mantle of mainly small condensed molecules. Non-thermal desorption plays a crucial role in the exchanges between these grains and the gas phase. Condensed small molecules also play a role in other regions of space, such as the variety of icy objects of the solar system (comets, moons, dwarf planets and others), and therefore the planetary science aspect of photon and electron irradiation of molecular ices must not be neglected.
The interstellar medium
The atomic constitution of the interstellar medium (ISM) is, by mass, 70 % hydrogen, 28 % helium and 2 % other heavier atoms. 99 % of matter is made of gas, and 1 % of dust grains, which will be evoked again later. By number, the constitution becomes 91% hydrogen, 8.9 % helium and 0.1 % heavier atoms. In table I.1 the (solar) abundances relative to hydrogen of a few relevant atoms is reproduced from ref. . The number density of grains relative to hydrogen depends on the region but in those that are of interest for us, a typical value is 10−12.
By the end of 2018, 204 molecules had been detected in the interstellar medium , containing up to 13 atoms (excluding the fullerenes C60, C+60 and C70). Links to an updated database and a discussion of these detections can be found in the previous reference. These detections are made across the electromagnetic spectrum, either in absorption or emission, from UV-visible (electronic transitions) and near/mid-infrared (vibrational transitions) to radio (mostly rotational transitions), which itself covers a large range going from far-infrared to centimeter wavelengths. Radio-astronomy has been by far the biggest contributor in number of detected molecules, for reasons that will become clear later. These molecules are not evenly distributed in the ISM: in fact, most of the interstellar medium does not contain molecules.
The interstellar medium is composed of regions of very diﬀerent physical conditions, with most of it being either very hot (106 K) and tenuous (< 10−2 particles.cm−3) ionized gas, or warmer (∼ 10 000 K) and slightly denser (10−1 particles.cm−3) gas that can be either ionized or neutral. In these regions the conditions of temperature and radiation make it impossible for molecules to be stable. What interests us here is the fraction of the ISM that is denser (> 10 particles.cm−3) and colder (< 300 K), enough that molecules can start forming, the so-called interstellar clouds. This is a small fraction in volume (< 5%) but a significant one in mass (30-60%) of the ISM.
According to the classification of Snow & McCall , four types of interstellar clouds can be distinguished: diﬀuse atomic, diﬀuse molecular, translucent, or dense, that are placed on a continuum of increasing density (101 to > 104 particles.cm−3), decreasing visual extinction (Av = 0 to 10)a and decreasing temperature (100 to 10 K). This classi-fication relies also on a chemical distinction: the transition from diﬀuse atomic to diﬀuse molecular is the transition from dominantly atomic H to molecular H2, the transition from diﬀuse molecular to translucent is the transition from dominantly C+ to C, and the transition from translucent to dense is the transition from C to CO.b Diﬀuse clouds are limited to a relatively simple chemistry because harsh radiations and low densities do not favor the formation of molecules for a long time, but molecules are still detected in these regions (CH, OH, CO, C3, ArH+…). In fact, diﬀuse clouds are among the best understood and constrained regions because the chemistry is simple, and because it is transparent to radiation and UV or IR emission diagnoses give directly access to e.g. the H2 density, which is a crucial parameter. Grains play an important role in these regions, even when they are too hot for molecules to reside for a significantly long time: they act as catalysts for the formation of H2 . When dust temperatures are low enough (< 100 K) in these regions, some molecules like water can start forming on grains, so that non-thermal desorption is not irrelevant there. In fact, the irradiation conditions are such that grains are kept bare by photodesorption .
In dense regions, observations are less varied: UV spectroscopy is not possible except at the edges, because of too strong dust scattering and molecule absorption. IR spec-troscopy in absorption can be done, but only if a background emitter exists behind the cloud or by targeting embedded stellar objects. The main diagnosis for gas phase mo-lecules remains radio-astronomy, with the emission of molecules rotationally excited by collisions or radiation. This has important consequences for what is detected: molecules that lack a dipole moment are invisible for radio-astronomy, while there is a bias towards some types of molecules that have a particularly large one. Important homonuclear di-atomic molecules like N2 or O2, and even more importantly the main component of the cloud H2, entirely elude direct detection. The low density and temperature make the gas phase chemistry in these regions driven by kinetics rather than thermodynamic equilib-rium, and only barrierless two-body reactions can take place, most notably ion-neutral reactions. The resultant chemistry is particularly exotic, with the observation of very reactive radicals that would be hard to isolate on earth and species that hurt the common sense of someone who only knows undergrad chemistry like myself (e.g. the important H+3 molecule).
aThe visual extinction is the degree of attenuation of light in the visible range by a medium. The Av number is a log scale: for Av = 0 there is no attenuation and Av = 5 corresponds to a factor of 100 attenuation.
b Another more « practical » distinction made by astronomers from an observational point of view is that a line of sight (the solid angle being probed by the telescope) where ices are visible in the infrared is dense, while one where ices are absent is diﬀuse.
These dense clouds, sometimes also called molecular clouds or dark clouds, are also the hosts of smaller-scale objects that are related to the star and planetary formation processes. This is usually presented in the form of a « cycle of matter » in the ISM. The story goes so: in molecular clouds (whose typical lifetime is of the order of a few My), some denser clumps can appear. These clumps will tend to grow denser and denser by gravitational attraction, forming so-called dense cores or pre-stellar cores, with very high extinctions and very low temperatures (down to 7-8 K). Eventually the gravitational col-lapse will form a protostellar object at the center and the temperature will start increasing, warming up the surrounding matter which is called a protostellar envelope. Where and when the temperature reaches the point where ice mantles sublimate, we reach a « hot core » (for massive star formation) or « hot corino » (for low mass star formation) phase which is very rich of a new gas phase chemistry, enriched by the ice mantle material and the possibility of neutral-neutral reactions. Matter is ejected from the protostar in the form of jets, creating shocks in the protostellar envelope. For low mass solar type stars, eventually the envelope dissipates, leaving a circumstellar disk, also called protoplanetary disk, surrounding the young star. The process from the start of the collapse into a dense core to the protoplanetary disk surrounding the star takes a total of about a million year. In the protoplanetary disk, future planets will form, along with the other objects we typ-ically find in our solar system (asteroids, comets, moons…), over a time scale of about 10 My. The cycle closes with the end of the life of the star (billions of years later for a solar type star), which expands and leaves a diﬀuse cloud, redistributing matter once again. Massive stars have much shorter lives and a more complex evolution, and their end leads to much more violent events like supernovae, where elements heavier than Fe are formed.
Dust and ice mantles
Dust grains in the ISM lock up a large part of the heavy elements available. They are mostly either silicate grains or carbonated grains, with varied forms of silicates (porous or crystalline, pyroxene and olivine…) and carbon (amorphous, diamondoid, aromatic or aliphatic…). Grains form mostly at the end of the life of a star, in the stellar winds created by the expansion and explosion, which create conditions where stellar material can be accreted in the form of a solid . The carbon abundance in the star will regulate the composition of dust: if the C/O ratio is > 1 mostly carbonaceous grains will be created, while in C/O < 1 conditions carbon is locked up in the formation of CO and silicate grains form. In addition to Si or C, looking at table I.1 these grains should also contain significant amounts of metallic atoms Mg and Fe, along with other refractory material in trace amounts (Al, Na, etc). Only part of these atoms are observed in the gas phase, and the rest is therefore presumed to be locked up in the grains. Models of dust include this material e.g. as nano-inclusions in silicate grains . Grains then evolve with time depending on the conditions they are exposed to: irradiation by UV photons and cosmic rays, or shocks, influence their nature and composition. Models of layered grains are often considered, with for example a core made of amorphous carbon or silicate and a « mantle » of aromatic-rich carbon .
The composition of dust can be observed by infrared absorption bands characteristic of either silicates or diﬀerent types of carbon. Also characteristic of dust is a continuous extinction curve going from the deep UV to the near IR, with a few famous features like the 217 nm bump and the diﬀuse interstellar bands (DIBs) in the visible/near-IR. The counterpart is an emission curve in the mid/far-IR. Models of dust attempt to reproduce correctly the features of these extinction and emission curves, constraining parameters such as the size distribution, shape and composition of the grains. In the ISM grains have sizes varying from about 1 nm to 1 µm, with a peak around 100 nm. The smaller grains are therefore basically macromolecules, with probably big PAHs (polycyclic aromatic hy-drocarbons) and fullerenes (C60, C70). Grains are also not particularly spherical: instead evidence from polarization of light by dust grains suggest more asymmetric shapes. In protoplanetary disks, dust grains can start to accrete and grow up to mm sizes. After reaching this size, further growth is diﬃcult and requires other and not yet completely un-derstood mechanisms to explain the formation of planetesimals . See the introduction chapters of  for more details on grains.
In the cold regions that are of interest for us, dust grains are covered by an icy mantle of molecules that can go up to a few hundred monolayers in thickness. These molecules either accrete from the gas phase or are formed directly on the grain through surface processes, and stay there when the grain temperature is lower than their sublimation temperature. Only infrared spectroscopy in absorption allows to probe ices, as condensed molecules do not have rotational transitions. As mentioned before, observations are made using either background stars behind the cloud or embedded stellar objects as infrared sources against which absorption features of condensed molecules can be seen. The observations that have been made of icy mantles are reviewed in Boogert .
The main species that have been detected in ice mantles are summarized in the graph of fig. I.1a, while an example of dust and ice features in the infrared spectrum of a massive young stellar object is shown in fig. I.1b. Ice mantles are composed mainly of water, with also large amounts of CO2 and CO. Depending on the line of sights CH3OH can be a very abundant molecule or be completely absent. NH3 and CH4 are observed at the few percent range. A number of other molecules have been detected in some line of sights, but whether this is representative of a global composition of ice mantles is not clear yet. According to the classification of Boogert , H2CO, OCN− and OCS are « likely detected » while some other molecules likes HCOOH, SO2, NH+4 or CH3CHO are « possibly detected ». In addition, upper limits for a large number of molecules (including very large upper limits on the abundance of hardly detectable species like N2 and O2) have been obtained. Solid phase infrared bands are broad and the sources weak, so that it is diﬃcult to obtain an inventory of molecules as varied and complex as is obtained in the gas phase, although there is undoubtedly numerous other molecules residing in ice mantles, as a result of solid phase (photo)chemistry of the main components mentioned.
The details of the absorption bands also give more information on the exact molecular environment of the molecules. The bands of CO and CO2 in particular have been used to deduce the existence of diﬀerent components in the ice. The most common distinction that is made is between polar and apolar ices, with the polar environment being mostly a water matrix and the apolar one a CO environment. This led to a picture of ice mantles as being roughly separated into two phases, linked with the way the mantles build up. At an early stage for the cloud, when grains are still bare and atoms are still available in the gas phase, the accretion of atoms on grains will lead to the formation of molecules like H2O, CH4 and NH3 by hydrogenation, and CO2. At a later stage when temperature drops so that the abundant gas phase CO freezes out onto the grains, the CO-rich second layer forms, with CO2 also being present. Hydrogenation of CO leads to CH3OH, with H2CO as an intermediate. This picture is coherent with the molecular environment suggested by the shapes of the absorption bands, but also with observations made as a function of extinction (for example CH3OH appears when extinctions are very high, corresponding to the massive CO freeze-out).
How much is the composition of the ice mantles conserved through the processes of star and planet formation is an important question that is diﬃcult to answer. It seems that an important part of the water is conserved as is from the pre-stellar core to the protoplanetary disk , while some portion has been sublimated in the protostellar phase and recondensed afterwards. Objects of the outer stellar system may then inherit this ice.
Solar system ices
Many objects in the solar system host molecular ices while being « airless » bodies (with no atmosphere or a very tenuous one), i.e. conditions that are close to the ones we study in the laboratory and in the context of interstellar grains. In the inner solar system (before the asteroid belt), temperatures are usually high, but in some permanently shadowed regions, water ice can exist. Its presence has notably been shown in craters of the Moon . In the outer solar system, there are many more objects of interests: the icy moons of the giant planets, the trans-neptunian objects, and comets.
The well-studied icy moons of Jupiter are Callisto, Europa and Ganymede. Their surfaces are covered with water ice. The surface temperatures are too high to condense more volatile species like CO, N2 or CH4 but other molecules such as SO2 and CO2 have been detected . Around Saturn, Enceladus is the icy moon that has generated the most interest, and aside from the vast parts of the surface that are covered with nearly pure water ice, organics and CO2 have been detected in some places. The other moons, except Titan, are also airless icy bodies for the most part. More details can be found in Bennett et al. .
With the Neptunian satellite system, which is mostly its biggest moon Triton, we reach a point where more volatile molecules can be found in icy forms. Thus the surface of Triton is composed of N2, CH4, H2O, CO2, CO and C2H6, with N2 being a very dominant species. The very volatile N2, CO and CH4 are also detected on Pluto. On Charon, only water has been found, although evidence suggests the presence of processed organics at the poles originating from CH4 photolysis . Comets are the most interesting objects to compare to interstellar grains in terms of ice composition, because they are often believed to have conserved a « primitive » composition, representative of what was found in the circumstellar disk of the our Sun prior to planet formation. Indeed, the composition of cometary ices are well studied and the detailed review of Mumma et al.  shows that it is quite similar to the kind of composition found for interstellar ices, with some specific diﬀerences .
Table of contents :
I.1 Astrophysical context
I.1.1 Molecules in the universe
I.1.1.1 The interstellar medium
I.1.1.2 Dust and ice mantles
I.1.1.3 Solar system ices
I.1.1.4 Sources of irradiations
I.1.2 Role of non-thermal desorption
I.1.2.1 Different non-thermal desorption processes
I.1.2.2 Astrochemical models
I.2 Vacuum technology context
I.2.1 Non-thermal desorption in vacuum dynamics
I.2.2 Non-thermal desorption in the context of accelerators and the LHC
I.2.2.1 Brief description of the LHC
I.2.2.2 The different sources of non-thermal desorption in accelerators
II Fundamental mechanisms of photon and electron-induced desorption
II.1 Position of the problem
II.2 Interaction of photons and electrons with molecular ices
II.2.1 Ices/molecular solids
II.2.2 Electronic transitions
II.2.3 From free molecules to molecular solids
II.2.4 Electron-matter interaction
II.2.4.1 Stopping power
II.2.4.3 Penetration and energy deposition profiles
II.2.4.4 Low-energy (< 20 eV) electrons
II.2.5 X-ray photons: core excitations and Auger decay
II.2.5.1 EXAFS, Shape resonances
II.3 Historic models of photon- and electron-induced desorption
II.3.1 The MGR model
II.3.2 Non-MGR models
II.4 A case study: desorption mechanisms from rare-gas solids
II.4.1 Electronic excitations in RGS: excitons
II.4.2 Desorption mechanisms
II.5 Desorption from molecular ices
III Experimental methods
III.1.1 Ultra-high vacuum
III.1.2 Mass spectrometry
III.1.2.1 Mass filtering
III.1.2.2 Ionization source
III.1.2.4 Residual gas analysis
III.1.2.5 Kinetic energy analysis with an electrostatic deflector
III.1.3 Ice growth
III.1.4 Temperature-programmed desorption
III.1.5 Infrared spectroscopy
III.1.6 Electron yield
III.1.7 Calibration of EID and PID
III.1.7.1 Absolute calibration methods
III.1.7.2 Relative calibration methods
III.1.7.3 Fragments and reflected light
III.2 Electron-induced desorption studies at CERN
III.2.1 The Multisystem set-up
III.2.2 Measurement procedure
III.3 SPICES II at LERMA
III.4 Synchrotron-based experiments
III.4.1 Synchrotron light and its advantages
III.4.2 Experiments on the DESIRS beamline
III.4.3 Experiments on the SEXTANTS beamline
III.5 Development of a UV laser desorption and spectroscopy set-up in the lab
III.5.1 UV and VUV laser desorption
III.5.1.1 VUV generation
III.5.1.2 Practical implementation
III.5.2 REMPI spectroscopy
III.5.3 Desorption + REMPI set-up
IV VUV photon-induced desorption
IV.1 Pure ice systems
IV.1.1.1 Recent studies on CO photodesorption
IV.1.1.2 Thickness and deposition temperature dependence of CO photodesorption
IV.1.1.3 Photodesorption mechanisms
IV.1.2.1 Synchrotron wavelength-resolved study
IV.1.2.2 NO gas phase REMPI
IV.1.2.3 NO desorption + REMPI
IV.1.4.1 Water ice structure and electronic spectrum
IV.1.4.2 Photodesorption yields and mechanisms in the literature
IV.1.4.3 Experimental results from synchrotron study
IV.1.5.1 Photodesorption spectra and yields
IV.1.6 Other organic molecules
IV.1.6.1 Photodesorption from HCOOH ice
IV.1.6.2 Photodesorption from organic molecules
IV.1.7 Perspectives and limits for pure ices
IV.2 Indirect desorption: model layered ices
IV.2.1 CO-induced indirect desorption
IV.2.1.1 Single layers
IV.2.1.2 Multiple layers and other systems
IV.2.2 H2O-induced desorption
IV.2.2.1 Results on single and multilayers on H2O and D2O
IV.2.3 Other systems
IV.3 Implementation in astrochemical models spectral dependence
V Electron-induced desorption
V.1 Chemically inactive pure ices: N2 and Ar
V.1.1 Interpretation of EID yield curves
V.1.4 Ar mixed with impurities: effects of ice composition
V.2 CO, CO2, H2O and the role of chemistry
V.3 Relevance of the data to astrophysical and accelerator contexts
V.4 Conclusions and perspectives on electron-induced desorption
VI X-ray photon-induced desorption
VI.1 H2O X-ray photodesorption
VI.1.1 Ice absorption spectroscopy and structure
VI.1.2 Desorption of neutral species, and astrophysical relevance of X-ray photodesorption
VI.1.3 Desorption of ions
VI.1.3.1 H− desorption
VI.1.3.2 Oxygen fragments desorption
VI.1.3.3 H+ desorption
VI.2 CO X-ray photodesorption
VI.2.1 Effects of the irradiation: TEY evolution
VI.2.2 Desorption of neutral species
VI.2.3 Desorption of ions
VI.2.3.1 Ice charging and ageing
VI.2.3.2 Mass spectrum of cations
VI.2.3.3 Spectral signatures
VI.2.4 Discussion on X-ray induced photochemistry
VI.2.4.1 CO irradiation chemistry
VI.2.4.2 Comparison of different probes of chemistry
VI.2.5 Astrophysical yields
VI.3 Conclusions on X-ray photodesorption
Appendix A: Calibration values