# Organic chemistry in the first phases of Solar-type protostars

Cecilia Ceccarelli  Univ. Grenoble Alpes, CNRS, IPAG, 38000 Grenoble, France
Claudio Codella  INAF, Osservatorio Astrofisico di Arcetri, 50125 Firenze,
Italy Univ. Grenoble Alpes, CNRS, IPAG, 38000 Grenoble, France Nadia Balucani
Dip. di Chimica, Biologia e Biotecnologie, Università di Perugia, 06123
Perugia, Italy Dominique Bockelée-Morvan  LESIA, Obs. de Paris, PSL Research
University, CNRS, Sorbonne Univ., Univ. Paris Diderot, F-92195 Meudon, France
Eric Herbst  Dep. of Chemistry, University of Virginia, PO Box 400319,
Charlottesville, VA 22904, USA Charlotte Vastel  IRAP, Université de Toulouse,
CNRS, CNES, UPS, (Toulouse), France Paola Caselli Max-Planck-Institut für
extraterrestrische Physik (MPE), 85748 Garching, Germany Cécile Favre  Univ.
Grenoble Alpes, CNRS, IPAG, 38000 Grenoble, Francee Bertrand Lefloch  Univ.
Grenoble Alpes, CNRS, IPAG, 38000 Grenoble, France Karin Oberg Center for
Astrophysics — Harvard & Smithsonian 60 Garden St., Cambridge, MA 02138, USA
Satoshi Yamamoto  Department of Physics, The University of Tokyo, Tokyo
113-0033, Japan

###### Abstract

Planetary systems such as our own are formed after a long process where matter
condenses from diffuse clouds to stars, planets, asteroids, comets and
residual dust, undergoing dramatic changes in physical and chemical state in
less than a few million years. Several studies have shown that the chemical
composition during the early formation of a Solar-type planetary system is a
powerful diagnostic to track the history of the system itself. Among the
approximately 270 molecules so far detected in the ISM, the so-called
interstellar complex organic molecules (iCOMs) are of particular interest both
because of their evolutionary diagnostic power and because they might be
potential precursors of biomolecules, which are at the basis of terrestrial
life. This Chapter focuses on the evolution of organic molecules during the
early stages of a Solar-type planetary system, represented by the prestellar,
Class 0/I and protoplanetary disk phases, and compares them with what is
observed presently in Solar System comets. Our twofold goal is to review the
processes at the base of organic chemistry during Solar-type star formation
and, in addition, to possibly provide constraints on the early history of our
own planetary system.

##  1 INTRODUCTION: A HISTORICAL PERSPECTIVE

Since the discovery of the first diatomic molecule, CH, in the interstellar
medium (ISM) at the end of the 1930’s, new molecular discoveries have been
made steadily over time. Most of the ISM molecules have been discovered by the
use of telescopes in the radio, millimeter and far-infrared (FIR) wavelengths,
where molecules have rotational and ro-vibrational lines. As of the present,
the number of identified ISM molecules is about 270, of which about 40%
contain at least six atoms111https://cdms.astro.uni-
koeln.de/classic/molecules. Importantly, all these ”large” molecules possess
at least one carbon atom, a fact that already suggests a potential link
between interstellar chemistry and terrestrial life. Indeed, in the zoo of
detected species, they are considered special and are called ”interstellar
Complex Organic Molecules” (iCOMs or COMs: _Herbst and van Dishoeck_ , 2009;
_Ceccarelli et al._ , 2017). The lower case ”i” is meant to emphasize that the
adjective ”complex” only applies in the context of the ISM.

Table 1 lists the names and the chemical formulae of the most commonly found
iCOMs in Solar-type star-forming regions.

Since the first discovery of formamide (NH2CHO) by _Rubin et al._ (1971),
iCOMs have aroused the curiosity of astronomers. The discovery that iCOMs are
numerous and very abundant in massive star forming regions (_Blake et al._ ,
1987) started to represent a serious challenge for astrochemists (e.g.
_Charnley et al._ , 1992). Eventually, the discovery of numerous and abundant
iCOMs in Solar-type protostars (_Cazaux et al._ , 2003) suggested a potential
link between interstellar chemistry and the emergence of life on Earth.
Understanding how iCOMs are formed and whether they are passed on to nascent
planets, comets and asteroids then became a major topic of astronomy.

At present, two major iCOM formation paths have been evoked in the literature;
these are reviewed in Sect. 2. The first one, of which a first version was in
vogue until about 2005, postulates that during the cold and quiescent phases
of the molecular cloud and prestellar core, the interstellar dust grains are
enveloped by ice mantles which sublimate when a protostar appears at the
centre and warms up the surrounding dust (_Millar et al._ , 1991; _Charnley et
al._ , 1992). The mantle components are ejected into the gas-phase and undergo
reactions that form new and more complex molecules. However, in 2005, a
breakthrough occurred with the experiment that showed that the electron
recombination of protonated methanol breaks the species into smaller species
(_Geppert et al._ , 2006) rather than producing methanol, which was the
assumption of the epoch. In addition, in 2004, theoretical chemical
computations showed that an abundant and common iCOM, methyl formate
(HCOOCH3), could not be formed by any known gas-phase reaction (_Horn et al._
, 2004). These negative results proved to be the death of the theory for most
iCOMs.

Table 1: List of the most common (i.e. more often detected) iCOMs detected in
Solar-type star-forming regions together with their binding energies (BE) and
pre-exponential factors $\nu_{des}$, when available (see § 2.3). References
refer to the latter values.

Name | Formula | BE | $\nu_{des}$ | Ref.a
---|---|---|---|---
|  | (K) | ($\times 10^{12}$ s-1 |
Methanol | CH3OH | 3770–8618 | $\sim 3\times 10^{5}$ | 1, 2
Formic acidb | HCOOH | 5382–10559 | $\sim 10^{5}-10^{6}$ | 1
Acetaldehyde | CH3CHO | 2809–6038 | $\sim 1\times 10^{5}$ | 3
Methyl formate | HCOOCH3 | – | – | –
Glycolaldehyde | HCO(CH2)OH | – | – | –
Acetic acid | CH3COOH | – | – | –
Dimethyl ether | CH3OCH3 | – | – | –
Acetone | CH3COCH3 | – | – | –
Ethanol | CH3CH2OH | $\sim 7000$ | $\sim 4\times 10^{6}$ | 4
Propanal | CH3CH2CHO | – | – | –
Ethylene glycol | (CH2OH)2 | $\sim 7100$ | $\sim 10^{5}-10^{6}$ | 5
Methyl cyanide | CH3CN | 4745–7652 | $\sim 2\times 10^{5}$ | 1
Ethyl cyanide | CH3CH2CN | – | – | –
Cyanoacetylene | HC5N | – | – | –
Formamide | NH2CHO | 5793–10960 | $\sim 4\times 10^{6}$ | 1, 6

NOTES: a References: 1, _Ferrero et al._ (2020); 2, _Minissale et al._ (2022);
3, _Ferrero et al._ (2022); 4, Private communication from J. Enrique-Romero;
5, _Bianchi et al._ (2022c); 6, _Chaabouni et al._ (2018). b Formic acid is
included in the list, although it is not an iCOM by strict definition, because
it contains three heavy atoms and is an important organic molecule.

If iCOMs cannot form in the gas-phase, they must be formed on the grain
surfaces according to several astrochemists. After all, the most abundant
molecule, H2, is known to be a grain-surface product (e.g. _van de Hulst_ ,
1946; _Hollenbach and Salpeter_ , 1970; _Vidali_ , 2013). Following this idea,
_Garrod and Herbst_ (2006) developed a new theory, in which grain-surface
chemistry can produce iCOMs during the warm-up phase (see § 2 for more
detail). Several laboratory experiments starting in the 90s had already shown
that, if an ice with a composition similar to that observed in the ISM by IR
absorption observations (see the review by _Boogert et al._ , 2015) is
illuminated by UV photons (_Schutte et al._ , 1992; _Bernstein et al._ , 1995,
1999) or energetic particles (_Strazzulla_ , 1997; _Palumbo et al._ , 1999),
several very complex molecules are formed on the ice. The most recent
experiments show that even amino acids and nucleobases can be found after
irradiation by UV photons or energetic particle irradiation (e.g. _Muñoz Caro
et al._ , 2002; _de Marcellus et al._ , 2015; _Oba et al._ , 2019). Several
studies have been carried out to better constrain the pathways to iCOMs
induced by photolysis and radiolysis (see e.g. the review by _Öberg_ , 2016).
Finally, recent experimental studies have shown that iCOMs and even glycine
could be formed via non-energetic processes on CO-rich ices irradiated by
thermal H atoms (e.g. _Qasim et al._ , 2019; _Chuang et al._ , 2020; _Ioppolo
et al._ , 2021).

1.5 sec2-chemistry-scheme_v5.png

Figure 1: Sketch of the two major paths invoked in the literature for the
formation of iCOMs. Both predict that, during the cold molecular
cloud/prestellar core stages (§ 3), the interstellar grains are covered by icy
mantles whose constituents are the result of the hydrogenation of atoms and
simple molecules, such as CO, and the oxidation of CO into CO2 (§ 2.2). After
this first step, the two paths differ as follows.
Left panel: In the Gas + grain chemistry path, the mantle components are
either partially or completely injected into the gas-phase by thermal and non-
thermal desorption processes (§ 2.3), where they undergo gas-phase reactions,
which convert them into iCOMs (§ 2.4). This path can occur either in cold
regions (i.e. in prestellar cores (§ 3), outer protoplanetary disks (§ 5) and
molecular outflows (§ 6)) via non-thermal desorption (§ 2.3.2) of the mantles
or in warm regions (i.e. in hot corinos (§ 4) and inner protoplanetary disks
(§ 5)) via thermal sublimation (§ 2.3.1) of the mantles.
Right panel: In the Only grain chemistry path, there are two major pathways:
diffusive and non-diffusive iCOM formation (§ 2.5) and iCOM formation via
reactions involving ice water molecules (§ 2.6). In the first path, the icy
mantles are processed by UV (photolysis) and CR (radiolysis) irradiation while
being formed (§ 2.5.2). In the diffusive path, When the dust temperature
increases, radicals on the mantles become mobile and diffuse (§ 2.5.3). When
radicals meet on the grain surfaces, they combine forming iCOMs (§ 2.5.4).
Some radicals are also predicted to meet and react in non-diffusive processes
during the cold phase. In the second path, radicals landing onto the grain
surfaces from the gas-phase react with water molecules of the ice, forming
iCOMs (§ 2.6).

After more than 15 years, the ”grain-surface does all” theory is starting to
have its own problems. First, since 2012, observations have shown that iCOMs
are numerous and relatively abundant (see § 3) in cold ($\sim 10$ K)
prestellar cores, where a warm-up phase is absent. Therefore, these
observations pose a significant challenge to the original version of the
_Garrod and Herbst_ (2006) theory and to astrochemistry in general. At the
same time, new laboratory experiments have proved that some neutral-neutral
gas-phase reactions, previously thought to be highly inefficient at the low
temperatures of the ISM, actually become very fast, thanks to the quantum
tunneling of pre-reactive complexes (_Shannon et al._ , 2013). Finally, deep
searches of the chemical literature have demonstrated that important gas-phase
reactions are missing in the astrochemical networks (_Balucani et al._ , 2015)
and new QM calculations have been enriching the databases with new gas-phase
routes (_Barone et al._ , 2015; _Skouteris et al._ , 2018). The game is once
again afoot. Very likely, both grain-surface and gas-phase chemistry
collaborate to yield the zoo of iCOMs now known to be present in Solar-type
star forming regions.

In addition, whatever the formation route, in cold and dense regions, such as
in prestellar cores and protoplanetary disks, some mechanism has to be able to
remove iCOMs from the grain mantles and inject them into the gas-phase, where
they are detected. This is a challenge for our understanding of the iCOM
chemistry and the current models. Several non-thermal processes have been
proposed in the literature, (see § 2), which depend on several poorly known
parameters. A crucial one, common to almost all non-thermal processes, is the
binding energy of the iCOMs, namely the energy necessary to release the frozen
molecule into the gas. New experiments and theoretical calculations are also
finally providing information on these processes and parameters, which are
absolutely necessary for the interpretation of the astronomical observations.

In this Chapter, we review the state of the art, by first introducing the
basics about iCOM chemistry (§ 2) and then describing and discussing each step
of the early phases of Solar-type star formation: prestellar cores (§ 3),
protostars (§ 4), protoplanetary disks (§ 5), molecular outflows, which offer
a unique chemistry laboratory (§ 6), and comets, which permit us to forge the
link between the present Solar System and the inheritance of interstellar
organic chemistry (§ 7). A final section (§ 8) concludes the Chapter.

##  2 iCOM FORMATION: GAS-PHASE AND GRAIN-SURFACE CHEMISTRY

###  2.1 Overview

In the ISM, two major paths have been invoked in the literature for the
formation of iCOMs, as summarized in Fig. 1 and here named Gas + grain
chemistry and Only grain chemistry, respectively. The first step, common to
both theories, occurs during the cold molecular cloud / prestellar core phase,
with the formation of icy mantles covering the interstellar grains. Atoms and
small molecules remain glued to the grains when they hit them and, there, they
are hydrogenated and oxidised (e.g. _Tielens and Hagen_ , 1982). In this way,
water, CO2, methanol, ammonia and methane become the major components of the
icy mantles, as observed in the Near Infrared (NIR) (e.g. _Boogert et al._ ,
2015). However, not all species frozen onto the grain surfaces become fully
hydrogenated so that some radicals are trapped into the mantle bulk (e.g.
_Taquet et al._ , 2012). After this first step, the two paths diverge.

In the Gas + grain chemistry path, depending on the temperature, thermal and
non-thermal processes release the grain mantle components into the gas-phase,
partially or completely. Once in the gas-phase, the mantle components react
and form, among other species, iCOMs other than methanol (e.g. _Vasyunin and
Herbst_ , 2013; _Balucani et al._ , 2015; _Skouteris et al._ , 2018). In the
Only grain chemistry path, the grain mantles are processed by UV (photolysis)
and cosmic-rays (CR) (radiolysis) irradiation, forming numerous radicals in
their bulk. These radicals become mobile when the dust temperature increases
to about 20–30 K and, when they meet, they react and form iCOMs (e.g. _Garrod
and Herbst_ , 2006; _Garrod et al._ , 2008; _Aikawa et al._ , 2020). In cold
($\leq 10-20$ K) environments, non-diffusive processes can also cause the
radicals to meet on the grain surface and react to form iCOMs (_Jin and
Garrod_ , 2020). Recently, the same non-diffusive mechanisms have also been
invoked to reproduce the iCOMs in warm regions (_Garrod et al._ , 2021), in
substitution of the diffusive models. In addition, radicals landing on the
grain surfaces can react with the water molecules of the icy mantles forming
iCOMs (e.g. _Rimola et al._ , 2018). Finally, the iCOMs formed on the grain
surfaces are released into the gas phase by thermal and non-thermal processes.

Probably, both gas-phase and grain-surface chemistry are at work, dominating
the formation of different iCOMs in different environments and times.

In the following, we will review the major processes which enter into the
formation of iCOMs as well as their presence in the gas-phase, where they are
observed: the grain mantle formation (§ 2.2) and desorption processes (§ 2.3),
the gas-phase (§ 2.4) and grain-surface (§ 2.5 and 2.6) chemistry, as well as
alternative theories (§ 2.7). The overview of all the above mentioned
processes and their interaction is schematised in Fig. 1, where it is also
reported the paragraphs where they are discussed. Our goal is to provide the
general ideas behind the astrochemical models that are used to interpret the
astronomical observations as well as the uncertainties caused by our limited
knowledge of the involved basic processes.

###  2.2 Mantle and methanol formation

The grain mantles are mostly made up of water, CO and CO2 with less abundant
hydrogenated species such as ammonia, formaldehyde and methanol, as shown by
IR solid-state astronomical observations (_Boogert et al._ , 2015). In
general, there is ample consensus on the fact that simple molecules (e.g. CO)
and atoms are hydrogenated and oxygenated when they freeze-out on the grain
surfaces. The simple iCOM methanol is believed to be prevalently formed in
this way, via the successive addition of H atoms to the frozen CO (e.g.
_Watanabe and Kouchi_ , 2002; _Rimola et al._ , 2014). However, even though
several laboratory experiments have been carried out and showed the
feasibility of the process, the details which enter into the astrochemical
models are not completely understood.

Two major questions are still open: (1) how the H atoms diffuse on the icy
grain surfaces and (2) what is the barrier of the CO + H and H2CO + H
reactions. For example, models often assume that H-atoms diffuse on the icy
grain surfaces via hopping (e.g. _Garrod_ , 2013), while recent sophisticated
quantum mechanical (QM) studies by _Senevirathne et al._ (2017) show that
below $\sim 15$ K tunneling dominates. Actually, the relative contribution of
H hopping versus tunneling is very difficult to quantify, since it depends on
the ice surface structure (_Hama et al._ , 2012). In addition, in the (many)
astrochemical models adopting the _Garrod_ (2013) recipes and parameters, the
diffusion energy of H is underestimated with respect to the theoretical
results by _Senevirathne et al._ (2017), with a consequent overestimate of the
H diffusion up to 5 orders of magnitude. Likewise, in the same models, the
energy barriers of the CO + H and H2CO + H reactions are assumed to be
rectangular with a guessed width to reproduce the observed iced methanol
abundance. However, when compared with the theoretical QM calculations by
_Song and Kästner_ (2017), the rates of CH3OH formation are underestimated by
up to 6 orders of magnitude. The result is that the astrochemical model
predictions of the frozen methanol abundance which use the _Garrod_ (2013)
recipe and parameters are fine-tuned to the 10 K molecular cloud case, but
they could be unreliable when considering different physical conditions (e.g.
densities and temperatures).

###  2.3 Mantle desorption

In interstellar clouds, molecules adsorbed onto mantles of icy material can
desorb either thermally or non-thermally.

####  2.3.1 Thermal desorption

The thermal desorption of a species bound to a substrate can be approximately
described by the Polanyi-Wigner equation (_Kolasinski_ , 2002). The first-
order solution pertains to adsorbate molecules that are not strongly
associated with each other. This occurs in most cases, when they are adsorbed
on a surface made of other species. Under these conditions, the solution for
the rate of desorption $k_{\rm des}$ (s-1) is given by:

$k_{\rm des}=\nu_{des}\exp({-E_{\rm des}/T})~{}=~{}\nu_{des}\exp({-BE/T})$ (1)

where $E_{\rm des}$ is the desorption energy, also known as the binding energy
(BE), given in units of temperature, and $\nu_{des}$ is the pre-exponential
factor, which depends on the adsorbate and surface (_Minissale et al._ ,
2022). We will discuss in some detail both quantities because they have an
important impact on the astrochemical models.

Thermal desorption is rarely important at 10 K except for atomic hydrogen and
a few other light physisorbed species where van der Waals and dispersion
forces keep the species bound to the surface because the desorption energy is
typically large enough to keep the adsorbate bound. As the temperature
increases, the more volatile adsorbates begin to desorb, and by a temperature
of typically 100–120 K, thermal desorption has become an important process for
a variety of molecular species.

An intermittent form of thermal desorption is powered by cosmic rays which,
when they strike the grains, heat them to a temperature that could cause the
sublimation of species frozen in the grain mantles (_Leger et al._ , 1985;
_Hasegawa and Herbst_ , 1993; _Roberts and Millar_ , 2007; _Sipilä et al._ ,
2021).

##### Overview of the BE and $\nu_{des}$ derivation

BE and $\nu_{des}$ can be either measured via laboratory experiments or
estimated via QM computations. _Minissale et al._ (2022) provide a recent and
detailed review of both techniques.

Figure 2: Sublimation temperature of frozen methanol (red) and formamide
(blue), as a function of the heating time, for a pre-exponential factor (Eq.
1) equal to $1\times 10^{12}$ (obtained using the _Hasegawa and Herbst_ (1993)
formula: solid line) and $3\times 10^{17}$ (obtained using the _Tait et al._
(2005) formula: dashed line) s-1, respectively (see text).

Here we briefly discuss their approach to make clear the limitations on the
published data used in the astrochemical models.

In laboratory experiments, BE and $\nu_{des}$ are simultaneously derived from
so-called TPD (Thermal Programmed Desorption) experiments. The species is
deposited or co-deposited over a cold (typically 10 K) surface, which is then
heated at a constant rate and the sublimation rate into the gas-phase measured
as a function of time. Therefore, TPD experiments strictly measure rates of
desorption. BE and $\nu_{des}$ are then derived from the desorption rate curve
by using Eq. 1 (e.g. _He et al._ , 2016).

Theoretical calculations compute the energy necessary for the adsorbate
species to leave the surface using the TST (Transition State Theory)
methodology. The theoretical works usually only provide the BE value, without
reporting $\nu_{des}$, even though it is a product of the calculations (e.g.
_Ferrero et al._ , 2022).

In general, the pre-exponential factor can be analytically calculated,
following the recipe given by _Tait et al._ (2005) (see the discussion in
_Minissale et al._ , 2022; _Ferrero et al._ , 2022). For relatively large
molecules, such as acetaldehyde, $\nu_{des}$ is $\sim 10^{18}$ s-1.

Both recent experiments and theoretical computations have shown that,
actually, there is not a unique BE for each species, but rather a distribution
of BEs that reflects that the bond of the adsorbed species depends on the site
and the species orientation (e.g. _Ferrero et al._ , 2020; _Bovolenta et al._
, 2020; _Tinacci et al._ , 2022b). Astrochemical models often, but not always
(e.g. _Grassi et al._ , 2020), use one single value for the BE, assuming
$\nu_{des}$ to be $\sim 10^{12}$ s-1, following the recipe by _Hasegawa and
Herbst_ (1993), regardless of the source of the BE evaluation. However, as is
the case for acetaldehyde, $\nu_{des}$ could be as large as $10^{18}$ s-1,
depending on the species and the method used to derive the BE (e.g. _Ferrero
et al._ , 2022). A review of the $\nu_{des}$ from their recipe is reported in
_Minissale et al._ (2022, see also their Fig. 14).

##### BE and $\nu_{des}$ of iCOMs

In general, only the BE and $\nu_{des}$ of a few species have been measured in
the laboratory (e.g. _Collings et al._ , 2004; _He et al._ , 2015, 2016;
_Corazzi et al._ , 2021; _Ferrero et al._ , 2022) or computed (e.g. _Wakelam
et al._ , 2017; _Das et al._ , 2018; _Ferrero et al._ , 2020, 2022; _Tinacci
et al._ , 2022b). The computations refer to the bond energy of one water
molecule (_Wakelam et al._ , 2017), a more complex situation with up to six
water molecules (_Das et al._ , 2018), and to models of amorphous water ice
clusters (_Ferrero et al._ , 2020, 2022; _Tinacci et al._ , 2022b), the most
reliable. The iCOMs where the BE and $\nu_{des}$ have been either measured or
evaluated on an ice cluster of more than 20 water molecules are listed in Tab.
1. Clearly, there is a massive lack of data in this respect and this is a
first serious warning for the astrochemical models. To illustrate the impact
of the assumed pre-exponential factor value $\nu_{des}$ (Eq. 1), Figure 2
shows the sublimation temperature of frozen methanol and formamide as a
function of the heating time, which often, but not always, corresponds to the
age of the studied astronomical object. Two values are assumed for $\nu_{des}$
(see discussion above). To obtain the sublimation temperature we used the
half-life time (=ln(2)/$k_{\rm des}$) of the frozen species. Taking the
example of methanol, if the heating time is $10^{3}$ yr then the sublimation
temperature would be $\sim 120$ K, while if it is $10^{6}$ yr it is $\sim 105$
K. The difference due to the pre-exponential factor is $\sim$20–30 K.

In the context of iCOM formation on the grain surfaces by the combination of
radicals (HCO, CH3, NH2 et cetera), the BE of the latter is also a crucial
parameter because it defines the range within which the radicals remain on the
surface and can, therefore, react with each other. Actually, the diffusion
energy, which regulates whether radicals can move on the grain surfaces, is
also a function of BE. At the time of this Chapter writing, there are no
laboratory measurements of the binding energies of radicals, because they are
extremely difficult to obtain. Theoretical calculations on models of amorphous
water ice exist for a few molecular radicals: O2, OH, HCO, CH3 and NH2.

####  2.3.2 Non-Thermal desorption

There are a number of different mechanisms for non-thermal desorption at low
temperatures, only some of which are reasonably well studied. Common to all of
these mechanisms is a source of energy to augment the low temperature in the
grain mantle. Several mechanisms are discussed here.

##### Chemical or reactive desorption

Chemical or reactive desorption occurs when a surface reaction possesses
sufficient exothermicity to overcome the binding energy of the products. This
process is analogous to what chemists refer to as unimolecular dissociation,
which is governed by the well-known RRKM (Rice–Ramsperger–Kassel–Marcus)
theory, or its simpler version, the RRK model (e.g. _Di Giacomo_ , 2015). An
approximation to the fraction of products released into the gas-phase through
chemical desorption based on the RRK approximation predicts approximately 1%
(_Garrod et al._ , 2007).

More recent and more sophisticated studies appearing in the literature have
challenged this simplified and usual assumption. For example, _Pantaleone et
al._ (2020) have simulated the chemical desorption of the reaction H + CO
$\rightarrow$ HCO via ab initio molecular dynamics simulations. These authors
found that the reaction energy ($\sim$1.4 eV) is largely absorbed by the icy
grains in less than 1 ps, leaving the newly formed HCO radical with too little
energy to break the bonds with the water surface. As a consequence, no
chemical desorption is expected to occur in this case, as was also found in
the experimental studies by _Minissale et al._ (2016). On the other hand,
_Pantaleone et al._ (2021) studied the H + H $\rightarrow$ H2 reaction and
found that the reaction energy ($\sim$5 eV) is also largely absorbed by the
water ice but the remaining fraction ($\sim$2 eV) allows the newly formed H2
to be desorbed, thanks to its low BE. Therefore, each system behaves in a
different way and the result depends on the reaction energy, the efficiency of
the water ice to absorb it and the BE of the formed species. Finally two
systematic experimental studies have provided a semi-empirical theory, in
which the chemical desorption depends on the species and on the substrate
(_Minissale et al._ , 2016; _Yamamoto et al._ , 2019).

With respect to iCOMs, methanol is the only species for which estimates of the
chemical desorption efficiency exist (_Minissale et al._ , 2016; _Chuang et
al._ , 2018; _Pantaleone et al._ , 2020). Experiments showed that the chemical
desorption efficiency caused by the addition of H atoms to frozen CO and the
partially hydrogenated species HCO, H2CO, CH3O and CH3OH, is very low ($\leq$2
%) on water-rich surfaces, (_Minissale et al._ , 2016; _Chuang et al._ ,
2018), in agreement with theoretical predictions (_Pantaleone et al._ , 2020).
Chemical desorption for other iCOMs forming on the grain-surfaces are at
present unknown.

##### Photodesorption

Another non-thermal mechanism often invoked in the literature is
photodesorption via UV radiation. This process has been studied in the
laboratory although this direct process can be confused with
photodissociation, in which the smaller products are the species that desorb.
The direct process appears to occur with a yield of $\sim 10^{-2}$ for CO
_Fayolle et al._ (2011). In these experiments, the photodesorption spectrum
tends to mimic the gas-phase spectrum (_Fayolle et al._ , 2011), namely the
gas-phase absorption spectrum of CO as a function of frequency and the
desorption spectrum of CO from a surface look the same in frequency with the
desorption spectrum having greater linewidths. A lower average efficiency of
$\sim 10^{-5}$ for methanol is thought to occur because of the simultaneous
photodissociation of the species (_Bertin et al._ , 2016; _Martín-Doménech et
al._ , 2016).

##### CR-induced desorption

Recent experiments have shown that sputtering by heavy cosmic rays can produce
relatively large quantities of gaseous methanol, originally frozen on the
grain mantles (_Dartois et al._ , 2015, 2019). Please note that this process
is not to be confused with the CR-induced thermal desorption described in §
2.3.1, in which the cosmic rays serve only to increase temperature.

##### Sputtering

In molecular shocks, frozen species can be injected into the gas-phase because
of sputtering caused by the ion-neutral drift velocity (_Draine_ , 1995;
_Flower and Pineau des Forets_ , 1995). Several models exist in the literature
to predict the amount of sputtered species, especially when the latter is SiO,
a common shock tracer (_Flower and Pineau des Forets_ , 1995; _Caselli et al._
, 1997; _Schilke et al._ , 1997). In addition to the treatment of the physical
and dynamical structure of the shocked gas, the crucial parameter in these
models is the so-called threshold energy of each species, which is the energy
of the projectile needed to extract the species from the grain mantle. There
is more than one treatment to evaluate the threshold energy and, not
surprisingly, all depend on the BE of the species on the surface (e.g. _May et
al._ , 2000).

Figure 3: Schematic view of the potential energy versus the reaction
coordinate (upper panels) and rate constant versus the inverse of the
temperature (1/T) (lower panels) of the three major classes of bimolecular
gas-phase reactions. The rate coefficients of bimolecular reactions mostly
obey the Arrhenius equation, $k(T)~{}=~{}A~{}exp(-{\rm E_{a}}/T)$, where the
temperature dependence of the rate coefficient is represented by a constant
factor A (only slightly dependent on temperature) multiplied by an exponential
function bearing the dependence on the temperature $T$. If we plot log($k$) as
a function of inverse temperature for a reaction characterized by a positive
activation energy (${\rm E_{a}}$), the slope is given by –${\rm E_{a}}$ (lower
left panel). In the case of reactions with a null activation energy, the rate
coefficient is temperature independent or shows a mild negative dependence on
T (middle panels). Only recently, there has been experimental evidence for the
presence of a weak non-covalent pre-reactive bound complex before the barrier
causes a bimodal trend with $T$: at high temperature, the system experiences
the presence of the barrier and an increase of $T$ corresponds to an increase
in the value of $k$. However, as the $T$ goes below a certain value, the
increased lifetime of the pre-reactive complex will favor the reaction via
quantum tunnelling through the activation energy barrier and the rate
coefficient has a positive trend with decreasing $T$.

Among iCOMs, only methanol has been the focus of specific shock studies (_May
et al._ , 2000; _Jiménez-Serra et al._ , 2008) even though it is predicted by
the Paris-Durham shock model, which is publicly
available222https://ism.obspm.fr/shock.html (_Lesaffre et al._ , 2013). While
there are studies on the threshold of species such as SiO and CO on silicate
surfaces (e.g. _May et al._ , 2000), no specific experimental or theoretical
studies exist about methanol on water- or CO- rich ices, to the best of our
knowledge. The methanol threshold energy is assumed to be equal to that of
water in the Paris-Durham shock model.

###  2.4 Gas-phase formation of iCOMs

####  2.4.1 Overview

Once the precursor molecules formed in the icy mantle of interstellar grains
have desorbed, molecular complexity can increase further through gas-phase
reactions. In addition, the main destruction routes included in astrochemical
models occur in the gas-phase.

In the low densities of Solar-type star forming regions, the most important
gas-phase reactions that lead to an increase in chemical complexity are
bimolecular reactions. This is because three-body collisions are too rare
except in the innermost and densest regions of protoplanetary disks (§ 5). The
same is true also for their destruction routes. In addition, given the low
temperatures, only reactions which are exothermic are possible. Exothermicity,
however, is insufficient because bimolecular reactions will not take place
unless the reactants collide with a certain amount of energy called the
”activation energy” (${\rm E_{a}}$) of the reaction. The activation energy can
be seen as the energy involved in breaking or weakening the original bonds
before the new ones are formed and only those collisions with energies equal
to or greater than the activation energy will result in a reaction. At low
temperatures, sometimes, tunneling under the activation energy could be
important. For this reason, bimolecular reactions of interest in
astrochemistry always involve either a radical or an ion; that is, transient
species with unpaired electrons, the reactions of which are usually
characterized by a very small ($\leq 1$ kJ/mol) or null activation energy.

Two major classes of reactions can form iCOMs in the gas-phase: ion-neutral
and neutral-neutral reactions. In Fig. 3, the potential energy profiles of
typical bimolecular reactions (upper panels) and the corresponding Arrhenius
plots (that is, the variation of the rate coefficient with the temperature,
lower panels) are shown. If a reaction is characterized by an activation
energy, ${\rm E_{a}}$, the rate coefficient increases with increasing
temperature. If the activation energy is null (barrierless reactions), the
rate coefficient is either temperature independent or exhibits a slightly
negative dependence. For this reason, until recently only reactions with null
${\rm E_{a}}$ have been considered in astrochemical models. However, recent
experiments based on the CRESU technique have shown a very interesting effect:
the presence of a pre-reactive complex before the system surpasses the
activation energy can make the reaction efficient also at temperatures that
are too low to overcome the barrier, via the tunneling effect (see below).
Note that Fig. 3 refers mainly to neutral-neutral systems; the rate constants
of ion-neutral reactions have more restrictive possibilities. Reactions with
non-polar neutrals are temperature-independent, while reactions with polar
ions are inversely dependent on the square root of temperature.

We will briefly describe ion-neutral and neutral-neutral reactions in the
following subsections. Finally, §2.8 will discuss the state-of-art of
astrochemical gas-phase reactions networks.

####  2.4.2 Ion-neutral reactions

As briefly mentioned in the Introduction, interstellar gas-phase chemistry was
initially based mostly on ion-neutral reactions, because of their large rate
coefficients at the low temperatures of the ISM (_Herbst and Klemperer_ ,
1973; _Watson_ , 1974; _Dalgarno and Black_ , 1976). The final step in the
complex ion-neutral reaction chains was supposed to be the conversion of the
ionic species into their neutral counterparts, via electron-ion recombination
reactions with the release of one H atom and the preservation of the ion
molecular skeleton. Later on, however, it was shown that electron-ion
recombination reactions are strongly dissociative because of the large amount
of energy released in the process and that the major product channels most
often consists of small neutral moieties (_Geppert et al._ , 2006). Product
channels dominated by the release of one hydrogen atom, on the other hand,
rarely account for more than a few percent of products.

####  2.4.3 Neutral-neutral reactions

The role of neutral-neutral reactions was only recognized in the 90’s after
the introduction of an experimental method that allowed investigation of the
kinetics of bimolecular reactions at very low temperatures. This technique
possesses the acronym CRESU333Cinétique de Réaction en Ecoulement Supersonique
Uniforme, which stands for Reaction Kinetics in a Uniform Supersonic
Expansion. Note that the CRESU apparatus can also be used for ion-neutral
reactions. (_Smith and Rowe_ , 2000).

At present, the characterization of neutral-neutral reactions relies on three
complementary approaches: (1) the CRESU technique, to derive low temperature
global rate coefficients; (2) collision-free experiments, to derive the nature
of the primary reaction products and their branching ratios; and (3) quantum
theory, to derive potentials and rates for chemical reactions. This combined
effort is necessary because the CRESU technique is able to characterize gas-
phase reactions under the appropriate temperature conditions, but cannot
reproduce the low number density environment typical of the ISM, while
collision-free beam experiments are able to reproduce the low number density
environment, but not the low temperature conditions, (e.g. _Casavecchia et
al._ , 2009, 2015), with the exception of the Bordeaux apparatus developed by
_Costes and Naulin_ (2010).

Finally, theoretical calculations are essential not only to support the
interpretation of the experimental data, but also to assist the extrapolation
of the experimental results. In addition, when no experimental data are
available, a theoretical characterization of a gas-phase reaction can provide
a realistic estimate of the reaction rate coefficients and product branching
ratios (e.g. _Sleiman et al._ , 2018; _Ocaña et al._ , 2018, 2019; _Nguyen et
al._ , 2019).

Reactions characterized by a pre-reactive complex and an activation energy
have recently attracted much attention (see Fig. 3). In particular, the
reactions of the abundant OH radical with several oxygenated iCOMs fall in
this category and their investigation has had a profound influence in the
understanding of interstellar chemistry. The effect is double: on the one
hand, the reaction of oxygenated iCOMS with OH must be considered as an
efficient destruction route of those species because of the abundance of the
OH radical. On the other hand, the radicals formed by the H-abstraction
process can react further and be involved in the formation of other iCOMs
(e.g. _Shannon et al._ , 2013; _Balucani et al._ , 2015; _Skouteris et al._ ,
2018).

The most important case of this family is the reaction between OH and CH3OH.
This reaction is known to have an experimental activation energy ($\sim 10$ kJ
mol-1) (e.g. _Shannon et al._ , 2013; _Roncero et al._ , 2018). In spite of
that, below 200 K the rate coefficient changes its slope with the temperature
and starts increasing with decreasing temperature (see Fig. 3). The $k$ values
below 100 K can be larger by a factor of 103 with respect to the room
temperature value. The explanation for such a strong effect has been
rationalized by invoking the role of a pre-reactive complex (_Shannon et al._
, 2013; _Roncero et al._ , 2018), which is formed by the strong H-bond
interaction between the OH radical and the -OH functional group of methanol.
At very low temperature, the pre-reactive complex, characterized by a
relatively deep potential well, can live for a time long enough to have a
significant probability to tunnel through the barrier rather than re-
dissociate back to reactants (_Heard_ , 2018).

Concerning the reactions for which no experimental data exist but theoretical
computations have been carried out, illustrative successful examples are the
reaction NH2 \+ H2CO (_Barone et al._ , 2015; _Skouteris et al._ , 2017), the
only formation gas-phase route of formamide known so far, and the reaction O +
C2H5O (_Skouteris et al._ , 2018) on which an entire reaction chain to form
very common iCOMs is derived, including glycolaldehyde and acetaldehyde.

Finally, barrierless reactions mostly occur via the formation of one (or more)
covalently bound intermediate(s) (see middle panel of Fig. 3). In the absence
of three-body collisions, the intermediate cannot be stabilized by dispersing
part of its internal energy and, if exothermic exit channels are open, its
high internal energy content leads to its fragmentation into one or more sets
of bimolecular products. In astrochemistry, however, radiative association
might occur, i.e., the intermediate can lose part of its internal energy by
emitting a photon (_Herbst_ , 1985). Normally, these processes are not
competitive because the typical reaction time is of the order of picosecond
while spontaneous emission occurs within milliseconds. However, if no
exothermic channels exist and the intermediate can only re-dissociate back to
reactants, radiative association can occur with a significant rate
coefficient. An important case is represented by the reaction CH3 \+ CH3O
forming dimethyl ether (_Vasyunin and Herbst_ , 2013; _Balucani et al._ ,
2015), for which theoretical computations predict large reaction constants
(_Tennis et al._ , 2021).

1.7 sec2-radrad-reactions.png

Figure 4: Scheme of the iCOMs proposed to form on grain-surfaces via the
combination of radicals present on icy granular mantles. The darker gray cells
indicate the presence of activation barriers, whose value is assumed to be
that reported in the last row. Adapted from _Garrod et al._ (2008).

###  2.5 Diffusive and non-diffusive grain-surface formation of iCOMs

####  2.5.1 Diffusive grain-surface chemistry

This theory, inspired by numerous laboratory experiments (e.g. _Öberg et al._
, 2009; _Linnartz et al._ , 2015; _Kaiser et al._ , 2015), was first proposed
by _Garrod and Herbst_ (2006) and is based on three fundamental steps: (i) the
formation of radicals by UV photons induced by CR irradiation (§ 2.5.2), (ii)
the diffusion of radicals when the dust temperature increases to 20–30 K, so
that radicals become mobile (§ 2.5.3), and (iii) the combination of radicals
into iCOMs (§ 2.5.4). We briefly describe the three steps in the following
subsections.

####  2.5.2 Mantle processing

During the formation of the grain mantles, which are enriched in water,
ammonia, methanol, formaldehyde and methane, CR-induced UV photons penetrate
the mantles and photo-dissociate the frozen species creating radicals. For
example, OH and CH3 are formed from the photo-dissociation of methanol. The
rate at which the radicals are formed is usually assumed to be equal to that
in the gas-phase, which is probably an overestimation. For example, based on
modeling of ice formation following the _Garrod et al._ (2008) scheme and
recipes, _Kalvāns_ (2018) estimated that in order to reproduce the observed
solid-phase carbon oxides and ammonia the rate of radical formation should be
a factor 3 lower than that in the gas-phase.

####  2.5.3 Radical diffusion

The motion of radicals involves hopping over potential barriers, the so-called
diffusion energy (DE) barriers, normally treated as a fraction of the
desorption barrier (_Hasegawa et al._ , 1992). Therefore, the diffusion of
radicals depends on their BE and the fraction DE/BE. We already noted that the
BEs of only a few radicals have been computed and published in the literature
(§ 2.3). To the best of our knowledge, there are no measurements nor
computations of the fraction DE/BE in radicals involved into the formation of
iCOMs. Various studies on non-radical species and atoms have reported DE/BE
ratios from 0.3 to 0.6 (_Karssemeijer and Cuppen_ , 2014; _Minissale et al._ ,
2016; _He et al._ , 2018). Finally, the coupled theoretical and experimental
study by _Ghesquière et al._ (2015) suggests that small molecules trapped in
water ice can diffuse at temperatures larger than $\sim$60 K driven by the
diffusion of the water molecules in the ice. However, at lower temperatures,
neither molecules or radicals seem to diffuse in the mantle bulk (_Theulé_ ,
2020) (against the majority of astrochemical models).

####  2.5.4 Radicals reactions

A crucial basic assumption of the Only grain chemistry path is that, when two
radicals meet on the grain surface, they react in a barrierless manner and
combine to form an iCOM (_Garrod et al._ , 2008). In this way, several iCOMs
are predicted to be formed, according to the scheme shown in Fig. 4. However,
recent QM computations have shown that radical-radical reactions on water ice
surfaces can have barriers and can also have competitive channels such as
H-abstraction processes that form simpler molecules. (e.g. _Enrique-Romero et
al._ , 2016; _Rimola et al._ , 2018; _Enrique-Romero et al._ , 2021, 2022). As
a consequence, the combination of radicals may not be the dominant channel or
even not occur at all on the interstellar grain surfaces. An emblematic and
well studied case is that of acetaldehyde, which is unlikely to form on the
grain surfaces via the association of HCO and CH3 according to QM simulations
(_Enrique-Romero et al._ , 2016, 2021),as opposed to that predicted by the
(_Garrod et al._ , 2008) scheme (Fig. 4). Experiments where radical diffusion
is very accurately controlled show that indeed this is the case: acetaldehyde
is not formed on the cold grain surfaces (_Gutiérrez-Quintanilla et al._ ,
2021) (though (_Martín-Doménech et al._ , 2020) did see it form, albeit at a
low efficiency). However, other formation reactions leading to species such as
ethane, methylamine, and ethylene glycol do not show barriers (_Enrique-Romero
et al._ , 2022). In conclusion, each system needs to be studied case by case
to understand what probably happens on the interstellar grain mantles.

####  2.5.5 Non-diffusive grain-surface chemistry

Laboratory experiments have shown that active chemistry also occurs on
granular surfaces via thermal, non-diffusive processes (e.g. _Theulé et al._ ,
2013), in particular in CO-rich ice surfaces (e.g. _Qasim et al._ , 2019;
_Ioppolo et al._ , 2021). Recently, based on earlier treatments using
stochastic simulations for the production of CO2 (_Chang and Herbst_ , 2014),
_Jin and Garrod_ (2020) developed a detailed theory of non-diffusive motion,
which can occur via a variety of mechanisms, the simplest of which is a
“three-body” process in which two species on the surface diffuse rapidly even
at low temperatures and form a two-body complex. If this complex is formed
adjacent to a reactive species, a rapid non-diffusive association can occur
since no diffusive barrier is involved. To transfer the mathematics from
stochastic simulations to standard kinetic treatments, one removes any
barriers to diffusive motion and replaces them by the probability that a
reactive third body lies adjacent to the complex and can react without any
chemical activation energy. The reaction is treated as instantaneous. Other
non-diffusive mechanisms include photodissociation leading to fragments, which
are formed adjacent to a reactive species. These and other non-diffusive
surface mechanisms have been included in a simulation of low-temperature
chemistry and found to form iCOMs more rapidly than via standard diffusive
chemistry (_Jin and Garrod_ , 2020). Moreover the approach can also be used
successfully to simulate low temperature laboratory experiments leading to
iCOMs, in particular glycine (_Ioppolo et al._ , 2021).

###  2.6 Radical-water ice reactions

Recent theoretical studies have shown that some iCOMs can be formed on water-
ice-rich surfaces by reactions of radicals, landing from the gas-phase into
the grain surface, with the water molecules of the ice itself. _Rimola et al._
(2018) studied the case of the formation of formamide by the reaction of CN
with H2O to form HNCOH, followed by an H atom addition. More recently,
_Perrero et al._ (2022) showed that ethanol can be formed by the reaction of
CCH with one water molecule of ice, plus successive hydrogenation by landing H
atoms. Finally, _Molpeceres et al._ (2021) showed that C atoms landing on the
icy grain surface react with water molecules and form formaldehyde. The study
of the radical-ice water route of formation is at its infancy and may give an
interesting alternative path to iCOM formation not involving diffusion on the
grain surfaces.

###  2.7 Other routes of iCOM formation

####  2.7.1 Rapid Radical Association

Another theory for the iCOM formation appeared in the literature is called
”Rapid Radical Association” (hereafter RRA) (_Rawlings et al._ , 2013). This
theory proposes that iCOMs are formed by three-body gas-phase reactions
between radicals in warm high density gas. This environment exists, for a very
short period of time, following the sudden and total sublimation of grain ice
mantles driven by the catastrophic recombination of trapped hydrogen atoms,
and other radicals, in the ice. These ”explosions” essentially cause short-
lived (100 ns) episodes of high density, high temperature gas which leads to
efficient three-body reactions.

####  2.7.2 Radiolysis

Another granular mechanism to form iCOMs invoked in the literature, especially
at low temperatures, is radiolysis, by which is meant the bombardment of icy
mantles by cosmic rays in interstellar clouds, or by high energy protons or
electrons in laboratory simulations (_Paulive et al._ , 2021). If the icy
mantles contain carbonaceous species such as CO or CO2, bombardment leads to
iCOM production via a complex sequences of reactions involving the formation
of ions and secondary electrons, which in turn form energetic radicals, which
lead to iCOMs. Although numerous experiments have been performed (e.g.
_Arumainayagam et al._ , 2019), it is only recently that theoretical
treatments for the low-temperature interstellar medium have been developed
(_Shingledecker et al._ , 2018) (see also _Reboussin et al._ , 2014). That
said, QM calculations providing an atomistic and precise view of the process
are not available in the literature.

###  2.8 State-of-the-art astrochemical reaction networks

Two major gas-phase networks are publicly available and used by astrochemical
models: KIDA444http://kida.obs.u-bordeaux1.fr/ (_Wakelam et al._ , 2015) and
UMIST555http://udfa.ajmarkwick.net/ (_McElroy et al._ , 2013) These networks
contain about 8000 reactions that involve about 500 species.

Unfortunately, as indicated in the two databases, the large majority of the
included reactions has never been studied in laboratory experiments or by
theoretical computations. Even in the cases of reactions that have been
investigated, the experimental conditions rarely reproduce those typical of
Solar-type star forming regions, either regarding the temperature or the
density. In the absence of laboratory data, rate coefficients and their
temperature dependence are mainly estimated with some chemical intuition or by
drawing analogies with similarly known processes. When the data are available
outside the temperature or density range of relevance, they are used as such
or are extrapolated.

However, both approaches can be seriously wrong because: 1) small details in
the molecular structure can induce a huge change in the chemical behaviour and
reasoning by analogy can cause severe mistakes; 2) the extrapolation of the
temperature dependence of rate coefficients outside the investigated range can
be very risky because a change in their reaction mechanism can alter the
temperature dependence in non-Arrhenius reactions, while a too high density
can lead to the stabilization of reaction intermediates or affect the reaction
outcome via secondary or three body collisions.

Writing down schemes of unknown reactions can also be dramatically wrong when
not even the enthalpy of formation of reactants and products is known. This is
particularly true when unstable species such as ions or radicals, or other
uncommon species are involved. For instance, _Tinacci et al._ (2022a) carried
out a systematic work to ”clean” the astrochemical reaction network GRETOBAPE-
GAS (https://aco-itn.oapd.inaf.it/aco-public-datasets/theoretical-chemistry-
calculations), which is based on the KIDA network, from the most obvious
source of error: the presence of endothermic reactions. Using the enthalpy of
formation derived by electronic structure calculations (_Woon and Herbst_ ,
2009; _Tinacci et al._ , 2021) about 17% of the entries calculated for
neutral-neutral reaction channels and 5% for ion-neutral reaction channels
resulted in endothermic reactions within the accuracy of the calculations.

Only one public database exists so far with a list of grain-surface reactions,
posted on the KIDA website (http://kida.obs.u-bordeaux1.fr/). The list is
based on the _Garrod et al._ (2008) and _Garrod_ (2013) reaction network and
presented in _Ruaud et al._ (2015).

1.7 sec3-PSCstructure.png

Figure 5: Sketch of the physical-chemical structure of the prototypical
prestellar core L1544 (see text and _Keto and Caselli_ (2010) for the
temperature and density profiles). X-axis report the distance from the center
of the prestellar core in au (bottom: assuming a distance of 140 pc) and
arcsec (top); Y axis report the volume density (left: blue curve) and gas/dust
temperature (right: red curve). The two vertical black lines show the location
where CO and NH2D start to catastrophically freeze-out onto the grain mantles,
disappearing from the gas-phase. They identify the radius of the CO and NH2D
depletion zones. The dust grains are illustrated as silicate cores covered
with the grain surface chemical composition (yellow) coated with the icy grain
mantle (in blue; thicker in the cold dense interiors of the core). The
temperature bump at $\sim$ 7000 au corresponds to the so-called photo-
desorption layer, a chemical zone between the FUV starlight pervading
interstellar space and the FUV produced by CR (see for example _Hollenbach et
al._ , 2009; _Caselli et al._ , 2012).

##  3 PRESTELLAR CORES: COLD DOES NOT MEAN POOR IN ORGANICS

###  3.1 Overview of the physical and chemical structure

The first step in the Solar-type star formation is represented by the starless
cores, dense condensations (_Benson and Myers_ , 1989) in the filaments of
molecular clouds (e.g. _André et al._ , 2019; _Könyves et al._ , 2015;
_Ladjelate et al._ , 2020; _Fiorellino et al._ , 2021). Prestellar cores
represent a particular sub-group of the starless cores. They are
gravitationally-bound objects (_André et al._ , 2014), with evidence of
contraction motions (_Keto et al._ , 2015; _Lee and Myers_ , 2011) caused by
the fact that matter slowly accumulates towards the center under gravitational
contraction counteracted by magnetic pressure and low levels of turbulence
(_McKee and Ostriker_ , 2007). Prestellar cores are relatively rare, as their
life times are short (_Könyves et al._ , 2015). The best studied ones are in
the Solar neighbourhood, at distances below 200 pc. Their sizes are about 104
au, with centrally concentrated structures (_Ward-Thompson et al._ , 1999),
reaching central visual extinctions larger than 50 mag and volume densities
larger than 105 cm-3 (_Crapsi et al._ , 2005; _Keto and Caselli_ , 2008),
where the temperature drops to about 7 K (e.g. _Crapsi et al._ , 2007; _Pagani
et al._ , 2007). They can show evidence of accretion of material from the
parental cloud (e.g. _Lee et al._ , 1999) as well as central contraction
(_Caselli et al._ , 2012).

In general, prestellar cores are ideal regions to study the initial steps of
chemistry, with no strong temperature gradient, no shocks, absence of internal
heating and, therefore, no outflows and protostellar feedback. The general
physical-chemical structure of the prototypical prestellar core, L1544, is
shown in the scheme of Fig. 5. The density increases towards the center while
the temperature decreases and reaches very low values (less than 10 K). This
physical structure defines three major zones from the chemical point of view:

  * 1

The molecular and iCOMs zone, where several molecules and, specifically, iCOMs
are observed.

  * 2

The depletion zone, where molecules such as CO freeze-out onto the grain
mantles and disappear from the gas-phase.

  * 3

The deuteration zone, where only H- and D- bearing species remain gaseous.

We will briefly describe the the three zones, starting from the first twos. We
will, on the contrary, discuss in details the organic chemical complexity in
the prestellar and starless cores, respectively.

###  3.2 Depletion and deuteration zones

In the central regions of prestellar cores, the large densities and low
temperatures lead to a zone where, with the exception of those containing only
H and D atoms, gaseous species freeze-out onto the cold grain surfaces,
forming thick icy mantles. In particular, CO depletion (e.g. _Caselli et al._
, 1999) goes hand in hand with a high deuterium fractionation, meaning a high
D/H ratio compared to the cosmic D/H value of $1.6\times 10^{-5}$ (_Linsky_ ,
2007), in the many detected deuterated molecules in the gas-phase (_Caselli et
al._ , 2002; _Bacmann et al._ , 2003; _Crapsi et al._ , 2005; _Ceccarelli et
al._ , 2007). This process can be explained by the fast reactions involving
the abundant ion H${}_{3}^{+}$, a product of the CR ionisation H2, with HD,
which is the main reservoir of D-atoms in cold molecular gas. The D-atoms are
then transferred to H2D+, D2H+ and even D${}_{3}^{+}$ (e.g. _Millar et al._ ,
1989; _Vastel et al._ , 2006), before being transferred to gas-phase species
(e.g. N2D+, DCO+ and NH2D) and, via dissociative recombination, to the D atoms
landing on the dust grain surfaces. The latter lead to an increase of the D/H
atomic ratio, by several orders of magnitude above the D cosmic abundance, and
an enhanced deuteration of the molecules synthesised on the grain surfaces
(e.g. H2CO, CH3OH, H2S and H2O). For example, the deuterated methanol, CH2DOH,
has been detected and mapped toward the prototypical prestellar core L1544
(_Bizzocchi et al._ , 2014; _Chacón-Tanarro et al._ , 2019). On the contrary,
H2D+ has been detected in only ten prestellar cores (_Caselli et al._ , 2003,
2008) and D2H+ in even less objects (_Vastel et al._ , 2004; _Parise et al._ ,
2011; _Harju et al._ , 2017; _Brünken et al._ , 2014), as these observations
are particularly difficult.

While CO depletion is important in the central 7000 au region of L1544, some
N-bearing species such as N2H+ and NH3 seem to trace the core center defined
by the mm and sub-mm observations of the dust continuum emission (e.g.
_Tafalla et al._ , 2004; _Spezzano et al._ , 2017). This was unexpected since
early measurements of the binding energy (BE) of N2 and CO provided similar
values (_Bisschop et al._ , 2006). However, more recent laboratory
measurements (_Fayolle et al._ , 2016; _Nguyen et al._ , 2018) and theoretical
QM calculations (_Ferrero et al._ , 2020) show that the two species have a
distribution of BEs which could possibly explain the observed segregation. In
less dense regions, the different timescales for the formation of CO versus N2
and NH3 (which also is formed by N2 in the gas-phase), about a factor ten,
could also explain the observed difference in the CO and NH3 emission. In
agreement with these new estimates of the N2 and CO BEs, high-J transitions
observations of N2D+ have provided evidence of depletion of this molecular ion
toward a prestellar core center (_Redaelli et al._ , 2019), without resolving
the depletion region. More recently, with the high sensitivity and spatial
resolution obtained with ALMA, the presence of a complete depletion zone,
where also gaseous ammonia disappears, has been discovered towards the L1544
within a 1700 au radius (_Caselli et al._ , 2022), as well as towards H-MM1
(_Pineda et al._ , 2022). This implies that most ($\sim$99.99%) species
heavier than He are locked onto the dust-grain ice mantles (_Caselli et al._ ,
2022).

###  3.3 Organic chemistry and iCOMs in L1544

As mentioned above, one of the breakthrough novelties on the organic chemistry
since the publication of the Protostars & Planets VI volume, is the iCOMs
detection in prestellar cores. As we will discuss later, the presence of
gaseous iCOMs in these cold objects poses a great challenge to the
astrochemical models for various reasons and, consequently, provides strong
constraints on how these molecules are formed and why they are present in the
gas-phase. Here we will review the observations, starting with the case of
L1544, which is the best studied prestellar core, then discussing the other
prestellar and starless cores in the next subsections.

As already hinted at above, among the known prestellar cores, L1544 stands
out. It is one of the densest and most dynamically evolved ones, towards which
H2D+ and H2O, for example, were detected for the first time in the cold ISM
(_Caselli et al._ , 2003, 2012). The physical and dynamical structure of L1544
has been computed by _Keto and Caselli_ (2010) and _Caselli et al._ (2012),
who predict a central temperature of $\sim$7 K, as actually measured by Crapsi
et al. (2007), and a peak number density $\sim$ 107 cm-3 (see Fig. 5).

The L1544 chemical structure has been the target of several observational
projects, in particular an unbiased spectral survey within the ASAI
(Astrochemical Surveys At IRAM) Large Program (https://www.oan.es/asai/
_Lefloch et al._ , 2018), aimed at exploring the chemical composition of the
earliest phases of Solar-type star forming regions. The survey covered the
70–110 GHz range and led to the detection of a very rich chemistry in L1544
and a better understanding of its molecular & iCOMs zone. For instance, more
than 20 S-bearing species have been detected in this survey and their modeling
led to the conclusions of a high elemental sulphur depletion in the gas phase
(_Vastel et al._ , 2018b; _Cernicharo et al._ , 2018). Similarly, many
N-bearing species have been detected which helped to better constrain the N
chemistry (_Vastel et al._ , 2015; _Quénard et al._ , 2017; _Vastel et al._ ,
2018a; _Hily-Blant et al._ , 2018; _Vastel et al._ , 2019).

More relevant to this review, several iCOMs have been detected towards L1544,
such as methanol (CH3OH), acetaldehyde (CH3CHO), dimethyl ether (CH3OCH3),
methyl cyanide (CH3CN) and vinyl cyanide (C2H3CN) (_Vastel et al._ , 2014;
_Bizzocchi et al._ , 2014; _Jiménez-Serra et al._ , 2016; _Vastel et al._ ,
2019). The study of methanol, which presents the brightest lines among the
detected iCOMs, is very instructive to understand the methanol spatial origin.
First, from the non-LTE analysis of the methanol lines, _Vastel et al._ (2014)
found that the bulk of the emission originates in the external ($\sim 8000$ au
from the center), relatively dense ($\sim 2\times 10^{4}$ cm-3) and cold
($\sim 10$ K) layers of L1544, what is defined as molecular and iCOMs zone in
Fig. 5. The average abundance of methanol in the external layer is about $\sim
6\times 10^{-9}$ with respect to H2. The maps of the methanol emission by
_Bizzocchi et al._ (2014) and _Spezzano et al._ (2016) and, at higher spatial
resolution, by _Punanova et al._ (2018) confirmed this interpretation and add
an important information on the methanol distribution. The methanol emission
is not uniformly distributed around the center of L1544, but it is peaked
towards the northeast, which is more shielded by the interstellar UV photons
than the southwest border of the L1544 condensation. On the same vein,
_Jiménez-Serra et al._ (2016) showed that the other O-bearing iCOMs abundance
is larger towards the methanol peak than towards its center by a factor 2–10
(within the uncertainties), again showing a definitive role of UV in the
presence of gaseous iCOMs. The abundance of iCOMs is about 10–100 times
smaller than that of methanol. Table 2 lists the measured abundance of iCOMs
in L1544.

Species | Column Density
---|---
| ($\times 10^{12}$ cm-2)
| Center | Methanol peak
CH3OH | 30 | –
CH3CHO | 0.5-1.2 | 3.2
HCOOCH3 | 4$\pm$4 | 2.3$\pm$1.4
CH3OCH3 | 1.5$\pm$0.2 | 0.77$\pm$0.16
t-HCOOH | 0.5 | –
CH2CHCN | 1.2$\pm$0.8 | 0.6$\pm$0.5
Table 2: Column densities of iCOMs detected towards L1544. The uncertainties,
if reported, are those quoted in the original works. References: _Vastel et
al._ (2004) and _Jiménez-Serra et al._ (2016).

While the northeast methanol peak is enriched with iCOMs, the southeast part
of the L1544 prestellar core is enriched in hydrocarbons and C-bearing chains.
Specifically, _Spezzano et al._ (2016) found that c-C2H3 abundance peaks
towards the south, where the condensation has a sharp edge and, consequently,
it is more exposed to the interstellar radiation field. Remarkably, in this
region, several carbon chains, such as C4H, C6H and C6H- are abundant (_Gupta
et al._ , 2009). The C-chain chemistry is so efficient that even long
cyanopolyynes with up to nine carbon atoms (HC9N) have been detected (_Bianchi
et al._ , 2022b).

The ensemble of these observations seems to provide a very clear message.
iCOMs and hydrocarbons/C-chains follow quite alternative routes for the
organic chemistry in prestellar cores. In the regions where iCOMs are
abundant, hydrocarbons/C-chains are scarce and vice versa. The presence of
strong UV illumination seems, at the moment, the major player in determining
the divergence. We will discuss the consequences of this finding in Sect. 4.

###  3.4 iCOMs in other prestellar cores

The presence of iCOMs is not a prerogative of L1544. Although difficult to
detect because their lines are weak, both because of the relatively low
abundance and gas temperature, a few iCOMs have been detected in several
prestellar cores. Notably, the survey by _Scibelli and Shirley_ (2020) towards
31 starless cores, of which many prestellar cores, in the Taurus Molecular
Cloud showed that methanol is present in all the targeted prestellar cores and
acetaldehyde in 70% of the sample666Please note that the survey contains also
prestellar cores.. Likewise, _Punanova et al._ (2022) mapped the methanol
emission in the Taurus L1495 filament and found emission associated with seven
starless cores. Therefore, iCOMs are likely a common product of the cold
chemistry.

As in L1544, the methanol emission in the _Scibelli and Shirley_ (2020)
sources originates in a cold (Tex=6–8 K) gas, and the derived column density
($\sim 10^{13}$ cm-2) is very similar to that found in L1544 ($\sim 3\times
10^{13}$ cm-2: _Vastel et al._ , 2014). Assuming that acetaldehyde emits in
the same gas than methanol, their CH3CHO/CH3OH is 0.02–0.2, with most sources
in the middle of the distribution (_Scibelli and Shirley_ , 2020)777However,
we emphasize that the error bars are relatively large (because of the few
lines detected).. The comparison with the CH3CHO/CH3OH$\sim 0.02$ in L1544
(_Vastel et al._ , 2014) would suggests a certain similarity among the
prestellar cores, when the uncertainties are taken into account, even though
this conclusion needs additional observations to rest on solid basis.

In a few additional prestellar cores, iCOMs other than methanol and
acetaldehyde have been detected. We will briefly describe their cases in the
following.

L1689B is a well-studied prestellar core in the $\rho$ Ophiuchus star-forming
region, which shows a moderate CO depletion (_Bacmann et al._ , 2002; _Jessop
and Ward-Thompson_ , 2001), probably sign of youth or a large CR/FUV
irradiation. L1689B is one of the first prestellar cores where three iCOMS
have been detected: acetaldehyde, dimethyl ether and methyl formate (_Bacmann
et al._ , 2012). The acetaldehyde column density is very similar to that found
in L1544 and the _Scibelli and Shirley_ (2020) cores, 1–3$\times 10^{13}$
cm-2. Likewise, dimethyl ether and methyl formate column densities are about
3–10 times lower than that of acetaldehyde.

H-MM1 is an intriguing prestellar core, located in the eastern side of the
L1688 cloud of the $\rho$ Ophiuchus complex, and illuminated on one side by a
strong UV field from nearby B-type stars. The asymmetry of the FUV field
illumination causes an enhancement of methanol emission in the shielded side
(_Harju et al._ , 2020), similarly to what happens in L1544 (see above).
Remarkably, also in H-MM1, as in L1544, ammonia is depleted towards its center
(_Pineda et al._ , 2022).

Finally, the first evidence of efficient production of iCOMs in cold
environments was reported by _Öberg et al._ (2010) and _Cernicharo et al._
(2012) toward B1-b, a dense core which hosts a Solar-type protostar. These
authors reported the detection of three iCOMs (methanol, methyl formate and
acetaldehyde) towards the quiescent core and not the outflow driven by the
protostar. However, the presence of the outflow could affect the observed
species, as also discussed by _Fuente et al._ (2016), who observed the same
source and found enhanced sulphur chemistry, typically an indication of the
presence of shocks (as sulphur is heavily depleted in starless cores; e.g.
_Ruffle et al._ , 1999).

###  3.5 iCOMs in starless cores

Starless cores which do not present evidence of infalling motion have a
relatively flat density profile and motions sometimes displaying oscillations
or even expansion (_Lada et al._ , 2003; _Tafalla et al._ , 2004). Thus, a
significant fraction of them could be transient and not evolve toward the
prestellar phase, which precedes the star formation (_Belloche et al._ ,
2011).

We have already mentioned above the study by _Scibelli and Shirley_ (2020),
which showed a high frequency of targeted starless cores presenting
acetaldehyde and all of them methanol emission. Here we discuss three cases.

L1521E is a starless core with no indications of infall motions and, on the
contrary, signs of a dynamical and chemical youth: modest central density
($10^{5}$ cm-3) and CO depletion ($\leq 2$). _Scibelli et al._ (2021) detected
acetaldehyde, dimethyl ether, methyl formate and acrylonitrile in addition to
methanol. The column densities of these iCOMs are about $10^{12}$ cm-12,
similar therefore to the prestellar cores, within the measurement
uncertainties. The youth of L1527E ($\leq 10^{5}$ yr) and the presence of
iCOMs at an early evolutionary stage testifies that the formation of the
latter is relatively fast, another constraint for the astrochemical models.

L1448 is also a starless core where _Jiménez-Serra et al._ (2021) found
interestingly relatively large abundances of N-bearing iCOMS but not O-bearing
ones, except for methanol.

Finally, by far TMC-1, in the Taurus Molecular Cloud complex, is the most
studied starless core, in particular the position known as ”cyanopolyyne
peak”. TMC-1 has a volume density similar to that in the external layer of
L1544 (at $\sim 10^{4}$ au), with n(H2) of a few $\times$104 cm-3. What made
it famous was the detection of abundant large cyanopolyynes, up to 11 carbons
(HC11N) (_Loomis et al._ , 2021). This has been interpreted by many authors as
an indication of a chemically young dense core (e.g. _Agúndez and Wakelam_ ,
2013), where carbon is not yet largely locked into CO and, therefore,
available for the observed rich carbon chemistry.

Indeed, TMC-1 is the site of many first molecule detections such as HC4NC
(_Cernicharo et al._ , 2020; _Xue et al._ , 2020), HC5NH+ (_Marcelino et al._
, 2020), H2CCCHC3N (_Shingledecker et al._ , 2021), indene (C9H8) (_Burkhardt
et al._ , 2021; _Cernicharo et al._ , 2021c), propargyl radical (CH2CCH)
(_Agúndez et al._ , 2021), propargyl cyanide (HCCCH2CN) (_McGuire et al._ ,
2020), allenyl acetylene (H2CCCHCCH) (_Cernicharo et al._ , 2021a), cyclic
aromatic molecules such as benzonitrile (c-C6H5CN) (_McGuire_ , 2018),
cyclopentadiene (c-C5H6) (_Cernicharo et al._ , 2021c) and
cyanocyclopentadiene (c-C5H5CN) (_McCarthy et al._ , 2021). Remarkably, the
detection of benzyne (ortho-C6H4) led _Cernicharo et al._ (2021b) to
hypothesise that, in TMC-1, PAHs are formed bottom-up from small carbon
chains, via gas-phase reactions of a few radicals with abundant gasoues
hydrocarbons, rather than top-down from carbonaceous-grain fragments.

###  3.6 Astrochemical models

As mentioned in the Introduction (§ 1), the discovery in 2012 of iCOMs in
prestellar cores (§ 3.4) was unexpected and unpredicted by the astrochemical
models. It still poses serious challenges to the current models for two major
reasons. The first one is the route of formation of iCOMs in the low
temperatures and high densities of the prestellar cores. The second challenge
is connected with the very presence of iCOMs in the gas-phase, as the BEs of
iCOMs are generally very large and close to that of water (e.g. _Ferrero et
al._ , 2020, 2022; _Corazzi et al._ , 2021, see also § 2).

In 2012, iCOMs formation models were prevalently based on the _Garrod and
Herbst_ (2006) scheme, described in Sect. 2 (see also Fig. 1). Briefly, iCOMs
were expected to form during the warm-up phase of the dust grains when the
protostar switches on and starts to heat the surrounding material (e.g. _Viti
and Williams_ , 1999; _Garrod et al._ , 2008), allowing mobility of the
radicals trapped within the icy mantles. These radicals could then react to
form iCOMs, which will then be released back in the gas phase when the dust
temperature exceeds the sublimation temperature of the specific species.
However, at the low temperatures found in prestellar cores, only hydrogen (and
deuterium) atoms can sweep across dust grain surfaces, thus allowing them to
react with other species. Depletion of carbon, oxygen and nitrogen atoms as
well as CO, as evidenced in prestellar cores (§ 3.2), leads to the production
of stable species such as methane, water, ammonia and methanol in the icy
mantles (see § 2). These species might be broken into radicals by the external
UV photons in the external shells of the core, or by the secondary UV photons
produced across the whole core by CR interacting with the abundant H2
molecules. The radicals which are not photodesorbed and remain in the icy
mantles are, however, too slow at these temperatures to effectively meet
together and, therefore, react to produce heavier iCOMs. Therefore, it appears
difficult for iCOMs other than methanol to be synthesized in the solid phase
in the cold prestellar cores via this scheme.

Two alternative routes have since been proposed (see § 2): the so-called ”non-
diffusive grain-surface chemistry” (_Jin and Garrod_ , 2020) (see below) and
the iCOMs formation in the gas-phase (_Vasyunin and Herbst_ , 2013; _Balucani
et al._ , 2015). In the latter scheme, iCOMs form in the gas-phase upon
release of methanol and other smaller iCOMs bricks from the grain surfaces. We
stress again the difficulty to having the iCOMs in the gas-phase, where they
are observed, for the first scheme and methanol in the second one, which
remains an open question (§ 2.3).

Detailed astrochemical models have been so far almost exclusively developed to
reproduce the iCOMs observed towards L1544 (_Vasyunin et al._ , 2017; _Goto et
al._ , 2021). These models use the core physical structure described by _Keto
and Caselli_ (2010) and _Chacón-Tanarro et al._ (2019), and start with the
production of methanol (CH3OH) on icy mantles rich in CO, at the edge of the
catastrophic freeze-out zone of prestellar cores (_Caselli et al._ , 1999).
Surface CO molecules are then efficiently hydrogenated into CH3OH, while the
CO-rich ice below the newly formed CH3OH molecule allows its efficient
reactive desorption (_Vasyunin et al._ , 2017) following the recipe by
_Minissale et al._ (2016) (but see the discussion in § 2). Finally, once
methanol and other radical/molecules are released into the gas phase, neutral-
neutral reactions can proceed to produce larger iCOMs (e.g. _Vasyunin and
Herbst_ , 2013; _Balucani et al._ , 2015; _Vasyunin et al._ , 2017). These
models predict that the peak of methanol and other iCOMs abundance is in a
layer at about 8000 au, in agreement with the observations (see § 3.3).
Besides, the predicted and observed abundances of gaseous methanol and iCOMs
agree within a factor 10, which is reasonable within the various
uncertainties. However, the predicted abundance of solid methanol is
significantly larger, by about a factor 4.5, than that observed at the edge of
the L1544 condensation by _Goto et al._ (2021).

Other models aiming at reproducing the iCOMs abundances observed towards
prestellar and, more generally, starless cores have also been published. The
model by _Jin and Garrod_ (2020) is based on the non-diffusive grain-surface
chemistry, described in § 2.5. This model, which uses a denser condensation
that the one used by _Goto et al._ (2021), approximately reproduces the
abundances of acetaldehyde and methyl formate observed towards L1544 at about
9000 au, but it fails to reproduce methyl formate by various orders of
magnitude. The model by _Holdship et al._ (2019a) is based on the Rapid
Radical Association, described in § 2.7. These authors tried to reproduce the
abundances of the iCOMs observed towards TMC-1 and concluded that the Rapid
Radical Association is not the dominant mechanism for iCOMs formation.

As a final note, the authors of the above models caution that ”a chemical
model is built on numerous parameters that we do not yet understand well”
(_Goto et al._ , 2021), both those regarding the microphysics of the involved
processes (e.g., efficiencies of grain-surface reactions, rate constants of
gas-phase reactions, binding energies, diffusion efficiencies on the surface:
Sect. 2) and the macroscopic ones, namely the profiles of the density, dust
and gas and temperatures, dust grain sizes, UV illumination, CR irradiation.
For example, in the initial work of _Vasyunin et al._ (2017), acetaldehyde was
produced via a gas-phase reaction with an impossibly large reaction rate
(_Vazart et al._ , 2020), and reduced in the successive works by _Jiménez-
Serra et al._ (2021) and _Scibelli et al._ (2021).

Much progress is needed on all these fronts and the next few years promise to
be crucial ones. For example, on the observational point of view, JWST is
expected to provide maps of methanol ice abundance in prestellar cores, like
that towards L1544 by _Goto et al._ (2021) but more sensitive, and perhaps
other iCOMs. Likewise, ALMA, NOEMA and VLA will provide more detections of
iCOMs and on a larger number of prestellar cores in different environments
allowing to better understand the influence of the environment on the early
organic chemistry.

1.5 sec4-protostars-structure.png

Figure 6: Physical and chemical structure of a typical Solar-type Class 0/I
protostar. It can be roughly divided into three zones, which are defined by
the status of the grain icy mantles formed during the prestellar phase (§3).
In Zone 1, the dust temperature is lower than the sublimation one of CO-rich
ices ($\sim 22$ K); in Zone 3 it is higher than the sublimation temperature of
H2O-rich ices ($\sim 100$ K); Zone 2 is the intermediate, lukewarm envelope
region. The upper x-axis reports the radius in arcsec, assuming a distance of
120 pc.

##  4 PROTOSTARS: THE RETAIL SHOPS OF ORGANICS

###  4.1 Overview of the physical and chemical structure of protostars

When the free collapse of matter towards the center starts, a protostar is
formed. Briefly, a protostar is made up of four main components: the central
object that will eventually become the new star, the infalling envelope from
which the central object accretes matter, the circumstellar disk in which
eventually planets and comets will form, and an outflow of ejected material.
The latter two components, the disk and the outflow, will be discussed in the
following Sects. 5 and 6. Here we focus on the infalling envelope and the
organic chemistry during the early phases of the protostar, when the observed
sources are called Class 0 and I.

As already discussed in the previous sections, the chemical structure and
evolution of the matter are tightly linked to its physical structure (e.g. the
density and temperature profiles) and the presence of ionising agents (e.g. UV
photons and CR). Given the presence of a central source of heat and centrally
peaked infalling envelope, the general physical-chemical structure of a
protostar can be divided into three major zones, illustrated in Fig. 6, from
the most external to the innermost ones, as follows:

  * 1

A cold envelope, characterised by a dust temperature lower than about 22 K,
where several molecules, including CO, are largely frozen into the grain
mantles, as in the prestellar cores described in Sect. 3.

  * 2

A lukewarm region, having a dust temperature between the CO and H2O
sublimation temperatures (approximately 22 and 100 K, respectively), where
organics can be present either in the form of carbon-chains or iCOMs or both.
In this region, iced methane sublimates: if abundant enough (see § 4.4), the
injection of methane gives rise to a zone enriched with unsaturated carbon-
chains, such as C4H or large cyanopolyynes. The sources where this zone is
dominating the chemical appearance are called WCCC (Warm Carbon Chain
Chemistry) sources (_Sakai and Yamamoto_ , 2013).

  * 3

A hot region, characterised by a dust temperature higher than about 100 K,
where the water-rich ices sublimate injecting into the gas-phase whatever
species is trapped in between. If frozen methanol and other relatively complex
species are abundant enough, they give rise to a zone very enriched with
iCOMs. This zone is called hot corino which, in its original definition, is a
compact source ($\leq 100$ au in radius) possessing high temperatures ($\geq
100$ K), high densities ($\geq 10^{7}$ cm-3) and showing the presence of iCOMs
in the region where the water-rich grain mantles sublimate (_Ceccarelli_ ,
2004). The typical spectrum of a hot corino is shown in Fig. 7.

WCCC sources and hot corinos are two chemically distinct classes of Solar-type
Class 0/I protostars, meaning that, observationally, their spectra are either
dominated by unsaturated carbon-chains or iCOMs. It is not by chance that the
first two distinct chemical classes of protostars, WCCC sources and hot
corinos, are based on which C-bearing species dominate their spectra, as
carbon is the element that allows the richest and most varied chemistry in
space (and on Earth!). At present, these two classes are the only ones
recognised to be present in Class 0/I protostars, even though there could
probably be a differentiation also with respect to N- and S-bearing molecules.
As often happens, hot corinos and WCCC sources are the two extremes, and
hybrid sources where the WCCC and hot corino chemistry appear on different
scales, as shown in Fig. 6, are found (e.g. _Imai et al._ , 2016; _Oya et al._
, 2017; _Higuchi et al._ , 2018; _Bouvier et al._ , 2020). As we will discuss
in §4.2, the number of known protostars with a well enough studied chemical
composition is very small at present, so that we are strongly biased in our
knowledge towards the hot corinos and WCCC sources. With the improving
observational facilities and the increasing number of observational surveys,
the situation promises to evolve in the next few years. For example, the ALMA
Large Program FAUST (Fifty AU STudy of the chemistry in the disk/envelope
system of Solar-type protostars; http://faust-alma.riken.jp, _Codella et al._
, 2021) is designed to reveal the chemical structure of several Class 0/I
sources.

2.0 sec4-hotcorino.png

Figure 7: Image (left panel) of the hot corino of the Class I protostar
NGC1333-SVS13 (Bianchi et al., in preparation) and the spectrum (right panel)
in the millimeter obtained by the NOEMA Large Program SOLIS (Seeds Of Life In
Space: _Ceccarelli et al._ , 2017).

###  4.2 The census of hot corinos and WCCC sources

##### Hot corinos

The first official detection of a hot corino, in IRAS16293-2422 (hereinafter
IRAS16293), dates back to 2003 (_Cazaux et al._ , 2003), although its
existence was suggested earlier by the observed formaldehyde abundance jump in
its inner region (_Ceccarelli et al._ , 2000) and the tentatively
identification of several iCOMs (_Ceccarelli et al._ , 2002). Later
interferometric observations showed that actually IRAS16293 harbors two hot
corinos, A and B, respectively (_Kuan et al._ , 2004; _Bottinelli et al._ ,
2004b; _Jørgensen et al._ , 2016). Soon after the first discovery, hot corinos
were also found towards NGC1333-IRAS4A (_Bottinelli et al._ , 2004a), NGC1333-
IRAS2A (_Jørgensen et al._ , 2005b) and NGC1333- IRAS4B (_Sakai et al._ ,
2006; _Bottinelli et al._ , 2007). Then for several years no new hot corino
was discovered, despite studies of formaldehyde and methanol emission had
suggested hot corinos also towards L1448-mm, L1448-N and L1157-mm (_Maret et
al._ , 2005; _Jørgensen et al._ , 2005a). Only in 2014, the number of hot
corinos with the detection of at least one iCOM other than methanol almost
doubled, adding to the list B1-a, SVS4-5 and B5-IRS1 (_Öberg et al._ , 2014).
These first seven hot corinos were all discovered using the single-dish
telescope IRAM-30m. The advent of ALMA and IRAM/NOEMA spurred new studies and,
finally, new surveys, such as CALYPSO (_Belloche et al._ , 2020), PEACHES
(_Yang et al._ , 2021) and ORANGES (_Bouvier et al._ , 2021, 2022) are
steadily increasing the number of known hot corinos. Another breakthrough was
obtained with the VLA observations at $\sim 25$ GHz of the binary system
NGC1333-IRAS4A, that showed that the dust is a major obstacle to the hot
corino detection (_De Simone et al._ , 2020a). At present, about two dozens
hot corinos are confirmed (meaning that they present line emission from more
than three iCOMs) and more are candidates (mostly when only methanol is
detected).

Before 2015, the fraction of known hot corinos in Class 0/I protostars was
$\geq$50%, which led to think that almost all Solar-type protostars went
trough the hot corino phase and that our Solar System may have not been an
exception (_Ceccarelli et al._ , 2007; _Öberg et al._ , 2014). However, the
number was heavily biased by the search criteria as, given the difficulty of
discovering hot corinos, the most probable sources were searched for. The new
surveys provide some contrasting results in terms of the fraction of hot
corinos in Class 0/I protostars: CALYPSO found $\sim 30$ % in a list of 26
protostars with a distance $\leq 420$ pc, when only sources with more than 3
detected iCOMs, referred in the article as ”confirmed” hot corinos are counted
(_Belloche et al._ , 2020), PEACHES found ($56\pm 14$)% in a large unbiased
sample of 50 protostars in Perseus (_Yang et al._ , 2021), and ORANGES found
($26\pm 23$)% in an unbiased sample of 19 protostars in Orion OMC-2/3
(_Bouvier et al._ , 2022), the last two surveys counting the sources with
detected methanol lines. When comparing the two unbiased samples, ORANGES and
PEACHES, which were carried out with comparable spatial resolutions (in au)
and sensitivities, there is a marginal, but potentially very interesting
evidence that the rate of hot corinos is affected by the environment.
Specifically, ORANGES studied the Orion region, where both high- and low- mass
stars are presently forming, whereas PEACHES studied the Perseus region, where
only low-mass protostars are present. In Orion, the presence of high-mass
forming stars makes the regions permeated by strong UV and CR irradiation
fields (e.g. _Ceccarelli et al._ , 2014b; _Favre et al._ , 2017), which may
reduce the onset of a complex organic chemistry leading to iCOMs in Solar-type
protostars (see § 4.4). Perseus, on the contrary, is not exposed to high UV/CR
irradiation even though it is known to be the theater of titanic clashes
between bubbles and clouds (_Dhabal et al._ , 2018; _De Simone et al._ ,
2022). Obviously, more regions need to be studied to draw firm conclusions on
whether environment affects, and at what extent, the chemical nature of the
forming Solar-type protostars.

##### WCCC sources

The discovery of the first WCCC protostar arrived in 2008 (_Sakai et al._ ,
2008), with the unexpected detection of abundant C4H2, C4H and CH3CCH towards
IRAS04368+2557 in L1527 and no iCOMs containing heavier atoms than C.
Subsequent observations showed that L1527 is enriched with a variety of
unsaturated carbon-chains, including cyanopolyynes with up to 9 carbons
(HC9N). As for the hot corinos, the early hunt for new WCCC sources ended up
in a frustrating addition of only two other protostars: B228 (_Sakai et al._ ,
2009) and TMC-1A (_Sakai et al._ , 2016) and a few other candidates.

Again, new surveys were carried out with single-dish telescopes, with the aim
of finding more WCCC sources and understanding the phenomenon, and ended up
with some contrasting results. _Graninger et al._ (2016) studied a sample of
15 Class 0/I protostars with measured (Spitzer) ice abundances and found the
co-existence and even correlation between C4H and CH3OH line emission,
suggesting that the observed protostars possess both a WCCC and a hot corino
zone. However, later, _Lindberg et al._ (2016) studied the same species
towards 16 protostars in the $\rho$ Ophiuchus molecular cloud complex and
found no correlation between C4H and CH3OH. On the contrary, they found a net
distinction between WCCC sources and (known) hot corinos. The extensive survey
of 36 sources in Perseus by _Higuchi et al._ (2018) used the CCH and CH3OH
line emission to discriminate between the two classes. Based on the two
prototypes of each chemical class, IRAS16293 and L1527, _Higuchi et al._
(2018) assumed that an abundance ratio [CCH]/[CH3OH]$\geq 2$ would define a
WCCC source, while [CCH]/[CH3OH]$\leq 0.5$ a hot corino. With this definition,
they found four new WCCC candidates. The most recent single-dish survey,
carried out by _Bouvier et al._ (2020), used again CCH and CH3OH line
emission, towards nine sources in the Orion OMC-2/3 clouds. All nine sources
present a [CCH]/[CH3OH]$\geq 2$, which would lead to conclude that they are
all WCCC sources. However, a careful analysis demonstrates that the single-
dish CCH and CH3OH line emission is likely dominated by the PDR (Photo-
Dissociation Region) enveloping the OMC-2/3 hosting clouds, making the
criterion flawed, when used in combination with single-dish observations only.

###  4.3 The census and abundances of detected iCOMs

Hot corinos are the Class 0/I protostars that present the richest chemical
composition, so that in this section we will focus on the iCOMs detected
there.

##### iCOM census and abundances

At present, more than twenty iCOMs have been detected in hot corinos, mostly
due to the spectacular spectral survey carried out with the ALMA
interferometer in the 300 GHz band by the PILS survey (Protostellar
Interferometric Line Survey: _Jørgensen et al._ , 2016). However, the vast
majority of iCOMs has only be detected towards IRAS16293 A and B (e.g.
_Jørgensen et al._ , 2016; _Ligterink et al._ , 2017; _Jørgensen et al._ ,
2018; _Calcutt et al._ , 2018; _Zeng et al._ , 2019; _Manigand et al._ , 2020,
2021), whereas only less than ten iCOMs have been detected in other hot
corinos (e.g. _Taquet et al._ , 2015; _Imai et al._ , 2016; _López-Sepulcre et
al._ , 2017; _Marcelino et al._ , 2018; _Bianchi et al._ , 2019; _Bergner et
al._ , 2019; _Belloche et al._ , 2020; _Ligterink et al._ , 2021; _Nazari et
al._ , 2021; _Martín-Doménech et al._ , 2021; _Chahine et al._ , 2022).

In general, measuring absolute abundances of iCOMs provides relatively
unreliable numbers, because of the uncertain H2 column density (for example,
the continuum does not necessarily trace the same gas where the iCOM emits)
and the problem of dust absorption of the iCOM lines (_De Simone et al._ ,
2020a). For example, _Zamponi et al._ (2021) found that the dust opacity
towards the IRAS16293 B peak exceeds 100 at 300 GHz, the frequency range
covered by the PILS survey. Consequently, even considering that the PILS
survey provided abundance one beam offset from the center, the absolute
abundances towards that source are likely substantially underestimated.

1.0 sec4-mf-dme-LC.png sec4-glyco-aceta.png

Figure 8: iCOMs in Class 0/I hot corinos. Top: Abundance ratio of dimethyl
ether over methyl formate (from _Chahine et al._ , 2022). Bottom: Column
densities of glycolaldehyde and acetaldehyde (_Vazart et al._ , 2020).

For the above reason, it is common use to plot relative iCOM abundance ratios
to understand how much chemistry is affected by the hot corino possible
different conditions and to provide constrains to the astrochemical models. As
an example, Fig. 8 shows the abundance ratio of dimethyl ether over methyl
formate in the known Class 0/I hot corinos (_Chahine et al._ , 2022). This
ratio, close to unity in all the known hot corinos, with the possible
exception of the hot corino HOPS108 in OMC-2 FIR4 (_Chahine et al._ , 2022),
suggests a mother-daughter or sister relation between these two species (e.g.
_Jaber et al._ , 2014; _Balucani et al._ , 2015), even though it is not a
proof (_Belloche et al._ , 2020). Another example is provided by the tight
correlation between the measured abundance of glycolaldehyde and acetaldehyde
in Class 0/I hot corinos, again suggesting a possible mother-daughter or
sister relation, where the two species may be one the mother of the other or
they both derive by a common mother, for example ethanol, as suggested by
_Vazart et al._ (2020). We will discuss both cases in § 4.4.

##### D fractionation in iCOMs

A recent and extremely important breakthrough has been the detection of
D-bearing iCOMs, other than methanol, where the D fractionation could be
measured (_Richard et al._ , 2013; _Coutens et al._ , 2016; _Jørgensen et al._
, 2018; _Persson et al._ , 2018; _Manigand et al._ , 2019; _Richard et al._ ,
2021; _Bianchi et al._ , 2022a): formamide, acetaldehyde, dimethyl ether,
methyl formate, glycolaldehyde, ethanol and ethyl cyanide. The detection and
abundance determination of deuterated iCOMs is important because it has the
great diagnostic power to discriminate the route of formation of the species,
as explained in § 4.4. In general, iCOMs show a deuterium fractionation (2–10
%) larger than that measured in water ($\leq 1$%) towards the same sources. As
already pointed out by _Ceccarelli et al._ (2014a), this systematic difference
is probably caused by the different formation time of water with respect to
iCOMs. Indeed, water is the first ice component to be formed whereas iCOMs are
formed after the depletion of CO onto the icy grain mantles, which enhances
the H2D+/H${}_{3}^{+}$ ratio in the gas-phase, at the base of the methanol and
formaldehyde deuteration.

##### iCOMs in the cold envelopes

Although with a much lower abundance, iCOMs are also present in the more
extended and cold envelopes of Class 0/I protostars, where the dust
temperature is lower than the icy mantle sublimation one. _Öberg et al._
(2010) found weak extended emission from iCOMs associated with the lukewarm
envelope of B1-b. Later, _Jaber et al._ (2014) found methyl formate, dimethyl
ether, acetaldehyde and formamide towards the cold envelope of IRAS16293.
Their abundances are about a factor 100 lower than those measured in the hot
corinos A+B, but still they are present. As mentioned in Sect. 3, iCOMS are
also found in prestellar cores, in the outer layers where the conditions are
similar to the lukewarm envelopes of the protostars, so that it is probable
that the same process is responsible for the presence of iCOMs in both
environments.

###  4.4 Astrochemical models

As the general ideas about the chemistry of iCOMs have been already discussed
in Sect. 2, here we only summarise some specific examples of few iCOMs.

##### Hot corino versus WCCC source chemistry

As discussed in § 4.2, Class 0/I protostars can be roughly classified into two
classes, hot corinos and WCCC sources, with respect to their organic contents.
Note, however, that this classification is not rigorous, because there are
many hybrid sources. What determines the chemical nature of a protostar is the
subject of debate. At present, two major reasons have been invoked in the
literature: a different time scale of the formation of the grain mantles
(_Sakai and Yamamoto_ , 2013), a different illumination from UV photons
(_Spezzano et al._ , 2016), or both. The above two hypothesis have in common
the idea that the differentiation is caused by the different major reservoir
of the frozen carbon: methanol versus methane. When methanol is more abundant,
it gives rise to the hot corino chemistry with the productions of iCOMs,
whereas when methane is more abundant the chemistry is dominated by the
formation of unsaturated carbon chains originating the WCCC sources.
Therefore, the above two hypothesis differ on what species dominates the
frozen carbon, methanol rather than methane. Methanol is thought to be formed
on the grain surfaces by the hydrogenation of CO, whereas methane by
hydrogenation of C. Therefore, the basic point is whether CO is abundant or
not on the grain surfaces with respect to C. The first possibility is that,
since methanol is formed on the grain surfaces after CO freezes-out at
relatively late times ($\geq 10^{5}$ yr), the formation time of the icy grain
mantles determines whether they are enriched with methanol rather than
methane. The second possibility is that CO is less abundant than C because of
the UV illumination that would photo-dissociate CO. Both mechanisms can
explain the existence of sources having the hybrid nature.

Models have been developed to quantitatively address this issue. _Aikawa et
al._ (2020) carried out a parameter study to understand what affects most the
appearance of a hot corino and WCCC source. They concluded that the chemical
appearance of the protostars depends on a complex function of time (i.e.
mantle formation duration), extinction (i.e. exposure to UV illumination), and
the thermal history (i.e. the dust temperature at the formation of the icy
mantles). On the other hand, _Kalvāns_ (2021) studied in detail the influence
of the UV illumination and CR irradiation, concluding that the WCCC sources
are the results of an enhancement of either of the twos.

##### Deuterated iCOMs

The iCOM deuteration follows a different path than the standard one initiated
by and completely dependent on the H2D+ ion at the moment of the molecule
formation (e.g. _Ceccarelli et al._ , 2014a). Indeed, since iCOMs are ”second-
generation” species their deuteration is inherited from their precursors (and
not directly from H${}_{3}^{+}$). For second generation here we mean that
iCOMs are not the products of hydrogenation, as the methanol, whose
deuteration therefore exclusively depend on the D/H atomic ratio on the grain
surfaces, which, in turn, depends on the H${}_{3}^{+}$/H2D+ ratio. For
example, acetaldehyde (CH3CHO) is thought to be formed by the combination of
HCO and CH3 on the grain-surfaces (_Garrod and Herbst_ , 2006) or by a series
of reactions starting from ethanol (CH3CH2OH) in the gas-phase (_Vazart et
al._ , 2020). Therefore, the deuteration of its two D-isomers, CH3CDO and
CH2DCHO, depends on the ratios of DCO/HCO and CH2D/CH3, if formed on the
grain-surfaces, and of CH2DCH2OH/CH3CH2OH, CH3CHDOH/CH3CH2OH and
CH3CH2OD/CH3CH2OH, if formed in the gas-phase.

In addition, the deuteration of the precursors is not just plainly passed to
the iCOM: the reaction leading to the iCOM itself will alter the deuteration,
because of the so-called Kinetic Isotope Effect (KIE), a well known effect in
chemistry. In practice, the final deuteration depends on the passed moiety and
how this is obtained, so that the iCOM deuteration depends on the deuteration
of its precursors and its formation route. If one has the possibility to
measure the deuteration of the iCOM and of its precursors, then this is
potentially an extremely powerful way to constrain the iCOM formation route.

A very nice example is provided by the case of formamide (NH2CHO), a species
that can in principle be formed both on the grain surfaces by the combination
of NH2 and HCO (_Garrod and Herbst_ , 2006) and in the gas-phase by the
reaction H2CO + NH2 (_Barone et al._ , 2015). In IRAS16293 B, the measured
abundances of the formamide deuterated isomers is in agreement with the
theoretical predictions of the gas-phase formation route (_Skouteris et al._ ,
2017), which may suggest that formamide is formed in the gas-phase in this
source. Unfortunately, no theoretical predictions exist so far on the
predicted deuteration in case of formamide formation on the grain surfaces so
that it is not possible to perform a similar significant comparison with the
observations. In the same vein, the deuterated forms of glycolaldehyde are
also in agreement with the theoretical predictions of its formation in the gas
(_Vazart et al._ , 2022), following the scheme starting from ethanol
(_Skouteris et al._ , 2018).

Figure 9: Sketch of the chemical structure of a generic protoplanetary disk.
In the hot surface layer, UV photons ionise atoms and photo-dissociate
molecules. In the warm molecular layer, where the UV photons are absorbed (by
the dust and gas), molecules, including iCOMs, can survive. In the cold
midplane, all species, except those containing only H and D atoms, freeze-out
onto the grain mantles.

##  5 PROTOPLANETARY DISKS: WHERE ORGANICS ARE BURIED

###  5.1 Overview of the physical and chemical structure of protostars

Circumstellar disks are expected to be present since the earliest phases of
the star forming process, i.e. the protostellar Class 0 phase described in the
previous section, and evolve into what are traditionally called protoplanetary
disks, in the Class I and II phases. Protoplanetary disks are complex
structures composed by a mixture of dust and gas and the chemistry occurring
in the gas-phase and/or on the grain surfaces reflects the diversity, the
inhomogeneity and the evolutionary stage of these objects.

Although their physical and chemical structure changes with time, from gas-
rich and small ($\sim 0.1~{}\mu$m) grains to gas-poor and pebbles, the
chemical structure can be approximately subdivided in three different zones,
shown in Fig. 9 (e.g. _Aikawa and Herbst_ , 1999; _Dullemond et al._ , 2007;
_Dutrey et al._ , 2014; _Walsh et al._ , 2014):

  * 1

The hot surface layer, where the chemistry is dominated by the UV photons. As
in the Photo-Dissociation Regions (PDRs), only ionised and neutral atoms are
present in this zone, because molecules are photodissociated.

  * 2

The warm molecular layer, where molecules are in the gas-phase. In this zone
the UV photons are largely absorbed but the dust temperature is large enough
for molecules to remain gaseous. This is the layer where iCOMs can be and are
detected (e.g. _Öberg et al._ , 2014).

  * 3

The cold (outer) midplane, where molecules freeze out onto dust grains and
only H and D containing species are in the gas-phase, such as H2D+
(_Ceccarelli and Dominik_ , 2005).

The freeze-out of different molecular species occurs at the disk radius and
height where the dust temperature is lower than the freeze-out temperature,
which in turn depends on the molecule BE. The loci of the gaseous-solid
transition state of a molecule is called snowline and it is specific to each
molecule. Although in the regions beyond the snowline molecules are
prevalently frozen into the grain mantles, a small fraction of them can be
released in the gas-phase thanks to the UV and X-rays photon illumination or
the CR irradiation (e.g. _Walsh et al._ , 2014; _Loomis et al._ , 2015), via
the non-thermal processes described in detail in 2.3.

In general, the detection of iCOMs in disks is very challenging because (i)
the dust along the line of sight might heavily absorb the line emission, and
(ii) the most chemically active regions are actually expected to have low
(beam-averaged) column densities because of the small vertical size of the
outer molecular layer and the small sizes of the warm molecular inner zone. As
a consequence, high-angular resolution and high sensitivities are crucial to
detect organic molecules (e.g. _Walsh et al._ , 2014). Remarkably,
protoplanetary disks show gaps in the dust continuum emission maps, which are
possible signatures of forming planets (e.g. _Fedele et al._ , 2016; _Sheehan
and Eisner_ , 2017, 2018; _Segura-Cox et al._ , 2020). Therefore, despite this
complex scenario and the observational challenge, it is essential to probe the
chemical composition of the young disks to understand whether more evolved
disks inherit the chemical complexity observed at the protostellar stage
(Sect. 4), as suggested by _Bianchi et al._ (2019) and _Drozdovskaya et al._
(2019).

###  5.2 iCOMs in Class 0 protoplanetary disks

Many recent observations show that circumstellar disks are present in Class 0
protostars, still embedded in infalling-rotating envelopes. The dust
absorption as well as the line emission from the infalling envelopes easily
mask the line emission from the disk itself, in particular weak iCOMs lines. A
remarkable exception is represented by the archetypal protostellar disk around
HH 212-mm, located in Orion B. The HH 212 system has been observed with ALMA
down to the spatial resolution of $\sim$ 10 au, resolving in dust continuum
emission the edge-on disk with a radius of about 60 au. The disk shows an
obscured equatorial plane caused by high optical depth near the midplane (_Lee
et al._ , 2017b). A large number of iCOMs with high abundances with respect to
H2 have been detected towards the HH 212 disk (see Fig. 10 from _Lee et al._ ,
2019a): CH3OH with $\sim 10^{-7}$, HCOOH with $\sim 10^{-9}$, CH3CHO with
$\sim 10^{-9}$, HCOOCH3 with $\sim 10^{-9}$ and NH2CHO with $\sim 10^{-10}$.
More specifically, iCOMs have been observed only in the vertical outer layers,
at $\pm$ 40 au from the equatorial plane. The kinematical analysis reveals
that the chemical enrichment occurs in rings rotating in the same way as both
the inner envelope (e.g. _Codella et al._ , 2018) and the SiO jet (_Lee et
al._ , 2017c). This picture raises the question of the origin of this iCOM-
rich structure, either in mild shocks associated with a centrifugal barrier or
passive heating from the protostar radiation or both (e.g. _Codella et al._ ,
2018, 2019; _Lee et al._ , 2017a; _Lee et al._ , 2019a). Similar results are
also reported for IRAS 16293-2422A (_Oya et al._ , 2016). Also, the lack of
iCOM emission towards the equatorial plane leaves the open question on a real
decrease of iCOM abundances or, again, a bias introduced by the dust emission,
which is optically thick at sub-millimeter frequencies. Observations at lower
frequencies, e.g. using the cm-spectral window, would be instructive in
revealing the chemical composition of the gas in the equatorial plane.

Figure 10: The protostellar disk associated with the Class 0 HH 212-mm object
(from _Lee et al._ , 2019a). iCOMs emission (contours) as imaged with ALMA
overlapped to the continuum map (color image). The line images have been
obtained by stacking a number of lines (as indicated in the parentheses) in
different transitions. The blue and red arrows indicate the axes of the blue-
and red-shifted jet components, respectively.

###  5.3 iCOMs in Class I/II protoplanetary disks

In the more evolved disks of Class I or Class II protostars, most of the
molecules are in the form of ices in the disk midplane and only a relatively
low column density remains gaseous. iCOM detections in Class I/II disks have
been rare, but the few detections that do exist provide vital information
about the iCOM reservoir available to forming planets, the evolution of iCOM
chemistry during star and planet formation, and the specific origins of iCOMs
in disks.

Up to date, few detections of iCOMs towards protoplanetary disks have been
reported. The detection of methyl cyanide (_Loomis et al._ , 2018), methanol
(_Walsh et al._ , 2016) and formaldehyde (_Öberg et al._ , 2017), which are
key species for the formation of other iCOMs, along with that of formic acid
(_Favre et al._ , 2018) towards the disk surrounding the very close (60 pc)
young Solar-type star T-Tau star TW Hya strongly suggests an active organic
chemistry taking place also at this stage (as discussed later) and/or non-
thermal desorption (see § 2.3) of previously formed organic molecules from the
grain mantles (see e.g. _Walsh et al._ , 2016). The recent survey of
protoplanetary disks MAPS (Molecules with ALMA at Planet-forming Scales;
_Oberg et al._ , 2021) shows that methyl cyanide, for example, is definitively
present in four out of the five targets (_Ilee et al._ , 2021). The emission
is concentrated in rings at distances that vary with the source or peaked
towards the center (see Fig. 11) and it originates in regions with
temperatures between 25 and 50 K, relatively close to the midplane, z/r $\sim$
0.1-0.2.

Figure 11: Continuum (first row) and molecular (bottom three rows) emission
distribution of the five Class II protoplanetary disks target of the ALMA
Large Project MAPS (_Oberg et al._ , 2021). The CH3CN emission (third row) is
distributed in rings or peaked towards the center and it is different from the
other imaged molecules HC3N (second row) or c-C3H2 (bottom row). Figure from
_Ilee et al._ (2021).

Methanol, a key iCOM, has been recently detected in the Class I disk IRAS
04302+2247 (_Podio et al._ , 2020a; _Garufi et al._ , 2021), and towards the
disk around the A-type star HD 100546 (_Booth et al._ , 2021). For sake of
completeness, methyl cyanide or formaldehyde (also crucial species for other
iCOMs formation) have been imaged towards a dozen of disks, such as e.g. TW
Hya, DG Tau B, DM Tau, HL Tau, MWC 480 (_Öberg et al._ , 2015; _Bergner et
al._ , 2018; _Loomis et al._ , 2018; _Podio et al._ , 2019; _Podio et al._ ,
2020b; _Pegues et al._ , 2020; _Garufi et al._ , 2021). Typical abundances
are: 10-12-10-10 (H2CO), 10-12-10-11 (CH3OH, HCOOH), and 10-13-10-12 (CH3CN).
The formaldehyde spatial distribution can be either centrally-peaked or
centrally-depressed showing an outer ring at a radius around 50 au (_Podio et
al._ , 2019; _Podio et al._ , 2020b; _Pegues et al._ , 2020; _Garufi et al._ ,
2021). The emission in the inner few au can indeed be masked by optically
depth effects. However, the detection of considerable formaldehyde emission
beyond the CO snowline shows that formaldehyde can be produced in the gas-
phase as well as through CO-regulated ice chemistry. Finally, the abundance of
methyl cyanide (and also that of HC3N) is relatively high compared to their
simpler relative, HCN (_Öberg et al._ , 2015; _Bergner et al._ , 2018).

In general, in protoplanetary Class I/II disks, complex nitriles are much
easier to detect than the most common iCOMs observed in the protostellar stage
(Sect. 4): methyl cyanide lines are considerably brighter towards the nearby
Class II disk around TW Hya compared to either methanol or formic acid
(_Loomis et al._ , 2018; _Walsh et al._ , 2016). This suggests, first, that,
at least in this disk, the iCOM inventory is not simply/completely inherited
from the prestellar or protostellar stages and, second, that the environmental
conditions in Class I/II disks favor the production of N-bearing iCOMs. This
latter conclusion is in agreement with the results from other studies focusing
on smaller molecules, which find that the disk warm molecular layers and disk
atmospheres likely have high C/O ratios (_Bergin et al._ , 2016; _Cleeves et
al._ , 2018; _Miotello et al._ , 2019; _Öberg and Bergin_ , 2021). The link
between the N- versus O- bearing iCOMs is discussed in § 5.5.

Finally, it is worth noting that the very recent detection of dimethyl ether
and, tentatively, methyl formate towards the Herbig IRS 48 transition disk by
_Brunken et al._ (2022) promises an era where it will be possible to discuss
the occurrence of iCOMs in protoplanetary disks and, consequently, the
chemistry at its basis.

###  5.4 iCOMs in Fu Ori stars protoplanetary disks

Protoplanetary disks surrounding FU Orionis stars (hereafter Fu Ori) are
precious laboratories where to study the disk chemical composition. Indeed,
the Fu Ori phenomenon leads the Solar-type star to vary the illumination of
the surrounding disk up to 2-3 orders of magnitude. The increased luminosity
causes the sublimation of the ices in larger portions of the disk making it
possible the observation of the molecular content just released from the
grains (see also § 7). The case of the disk surrounding the V883 Ori star is
an emblematic one. This protostar suggests a past outburst activity. A burst
was first announced in 1888 (but not measured: see _Connelley and Reipurth_ ,
2018), and later _Strom and Strom_ (1993) have measured a bolomotric
luminosity of $\sim$400 L⊙. The luminosity has been decreasing in the past
years to a present value of 218 L⊙ (_Fischer et al._ , 2019). V883 Ori is a
Class I protostar, where the surrounding infalling envelope is not yet
completely dissipated. Regarding the molecular content, five iCOMs (methanol,
13CH3OH, methyl cyanide, CH3CO, methyl formate) have been detected using ALMA
observations (_van ’t Hoff et al._ , 2018; _Lee et al._ , 2019b) (see also
Fig. 16). The relatively large emitting region
($\sim$0.4$\arcsec$–0.8$\arcsec$ in diameter equivalent to about 150–300 au)
coupled with the ALMA high spatial resolution permit to map the kinematics of
the gas emitting the iCOMs lines and demonstrate that, indeed, they originate
in a rotating disk, as shown in Fig. 12 (_Lee et al._ , 2019b). Interestingly
enough, the derived gas temperature, 100–120 K, and the measured molecules
column densities, 1016–1017 cm-2 (_Lee et al._ , 2019b), are similar to the
ones observed in typical hot corinos (_Ceccarelli_ , 2004).

Figure 12: The rotating disk around the Fu Ori object V833 Ori (from _Lee et
al._ , 2019b). Intensity weighted velocity images (colours) and integrated
intensity maps (contours) of the species reported in the upper-right corner.

###  5.5 Astrochemical models

The formation of iCOMs in disks have been explored theoretically, though the
number of specific disk chemistry model studies that focus on iCOMs has been
small, probably because of the few iCOM observations to date. One of the most
comprehensive study was carried out by _Walsh et al._ (2014), who predicted
that in situ disk chemistry can efficiently produce large quantities of iCOMs.
Whether these iCOMs are observable in the gas-phase depends on their non-
thermal sublimation efficiencies, however, which are poorly constrained for
iCOMs (see §2.3). In general, the desorption efficiencies are suspected to be
quite low, based on the experiments and computations on the paradigmatic iCOM,
CH3OH, (e.g. _Martín-Doménech et al._ , 2016; _Minissale et al._ , 2016;
_Chuang et al._ , 2018; _Pantaleone et al._ , 2020) and the relatively large
BE of the iCOMs. Whatever the mechanism of the iCOM formation, gas-phase (§
2.4) or grain-surface (§ 2.5)), and the non-thermal desorption efficiency (§
2.3), the vast majority of produced iCOMs end up frozen into the grain mantles
due to short freeze-out timescales.

Two recent models focused on the most commonly observed disk iCOM: CH3CN
(_Loomis et al._ , 2018; _Le Gal et al._ , 2020). Both models predict that in
situ grain-surface production could explain the observed abundances and also
reproduce the observed excitation temperature. In addition, _Le Gal et al._
(2020) and _Fedele and Favre_ (2020) also found that an enhanced C/O ratio
greatly facilitated CH3CN production, which may explain why this particular
iCOM is so abundant in disks. Needless to say, all the uncertainties on the
chemical processes incorporated in the astrochemical models and discussed in
§2 apply to the chemical disk models too, so that a word of caution is
mandatory on the interpretation of these theoretical results.

Do the model results in the previous paragraph imply that all observed disk
iCOMs are produced in the disk? Both models and observations show clear
evidence for substantial inheritance of iCOMs into disks from the preceding
evolutionary stages. Most of the time, however, this inherited reservoir would
be expected to be buried in the icy midplane and not be accessible to gas-
phase observations with e.g. ALMA. Recently, this scenario has been confirmed
by the CH3OH detection in two Herbig Ae disks that are themselves too warm to
host in situ CH3OH formation from CO ice hydrogenation (_Booth et al._ , 2021;
_van der Marel et al._ , 2021). Both disks belong to the class of the so-
called transition disks, with a large central cavity, which directly exposes
the icy disk interior to stellar radiation resulting in ice sublimation. The
picture that emerges is then one of both inheritance and in situ disk
chemistry. Molecular prestellar and protostellar iCOMs are incorporated into
Class II disks and constitute a rich chemical reservoir of O-rich organics.
Most of this reservoir is hidden from sight, in the icy disk midplane.
Meanwhile, in the (somewhat) warmer disk layers, the gas is C-rich, which
forms the basis for a gas and grain-surface O-poor iCOM chemistry, signified
by the presence of abundant complex nitriles. All together, this inherited and
in situ chemistry sets the stage for complex organic chemistry in forming
planetesimals, asteroids, comets and planets.

Table 3: List of iCOMs detected in protostellar outflow shocks and ordered
according to the presence of only O atoms (upper part of the table), N atoms
(middle table) and S atoms (bottom table). The iCOM maps were all obtained
towards the shocked site L1157-B1 and the outflows driven by NGC-1333 IRAS4 A1
and A2. Note that we did not include pure hydrocarbon species such as CH3CCH,
as done for the other sections of this review. Numbers in brackets refer to
the references in the Notes.

Element | Detected | Imaged
---|---|---
O | CH3OH [1], 13CH3OH [2,3], CH2DOH [4] | CH3OH [10], CH2DOH [11]
| CH3CHO [4], CH3OCH3 [7] | CH3CHO [12], CH3OCH3 [14]
| CH3CH2OH [5], HCOOCH3 [5], HCO(CH2)OH [7], HCOOH [5] |
N | NH2CHO [3], CH3CN [9], HC5N [8] | NH2CHO [13], CH3CN [9]
S | CH3SH [6] |

Notes: We give as reference the first detection reported in the literature
both for detection and interferometric imaging. [1] _Bachiller and Pérez
Gutiérrez_ (1997); [2] _Codella et al._ (2012)); [3] _Yamaguchi et al._
(2012); [4] _Sugimura et al._ (2011); [5] _Arce et al._ (2008) [6] _Lefloch et
al._ (2018); [7] _Lefloch et al._ (2017); [8] _Mendoza et al._ (2018); [9]
_Codella et al._ (2009); [10] _Benedettini et al._ (2007); [11] _Fontani et
al._ (2014); [12] _Codella et al._ (2015); [13] _Codella et al._ (2017); [14]
_De Simone et al._ (2020b).

##  6 MOLECULAR OUTFLOW SHOCKS: THE SPACE CHEMICAL LABORATORIES

###  6.1 The census of iCOMs in molecular shocks

The first observational evidence of iCOMs in protostellar outflows dates back
to the mid-1990s with the detection and the spectacular map of methanol in the
prototypical outflow of L1157 (and two other outflows) by _Bachiller et al._
(1995). More studies of methanol emission followed up towards a few Solar-type
Class 0 sources such as BHR71, NGC133-IRAS2 and confirmed the association of
bright CH3OH emission with shocked gas, related to non-steady mass-loss
ejections (_Bachiller and Pérez Gutiérrez_ , 1997; _Bachiller et al._ , 1998,
2001; _Jørgensen et al._ , 2004; _Parise et al._ , 2006). Multi-line radiative
transfer modelling of the CH3OH lines allowed to constrain the gas density and
temperature to be larger than the ambient gas, supporting the interpretation
of methanol being associated with shocked gas. The next significant step was
made by _Arce et al._ (2008), who reported the first detection, towards the B1
shock of the L1157 outflow, of other iCOMs than methanol: methyl formate,
formic acid, methyl cyanide and ethanol. Soon after, _Codella et al._ (2009)
reported the first high resolution map of an iCOM other than methanol, that of
the methyl cyanide, again towards L1157-B1, showing a distribution similar,
but not exactly the same, to that of the methanol one and definitively
different from that of SiO, another classical shock tracer. These results
definitely made the shock region L1157-B1 the prototype for studies on the
shock-driven chemical complexity.

Figure 13: Images of the L1157 southern outflow lobe, showing a chemical
segregation in L1157-B1. Left panel: The map of the CO (1-0) line emission
(white contours _Gueth et al._ , 1996) shows the morphology of the outflowing
gas. The (precessing) jet ejected by the central object L1157-mm excavated
several clumpy cavities, the brightest one being B1. The emission map of the
formamide line 41,4–-31,3 is shown in colour (_Codella et al._ , 2017). The
dashed circle indicates the primary beam of the NH2CHO image. The magenta and
white ellipses depict the synthesised beams of the CO (white) and formamide
(magenta), respectively. Right panels: The formamide spatial distribution is
compared with those of methanol (_Codella et al._ , 2020) and formaldehyde
(_Benedettini et al._ , 2013)(left), HDCO (_Fontani et al._ , 2014) and
acetaldehyde (_Codella et al._ , 2015, 2017) (middle), high-velocity SiO
(_Gueth et al._ , 1998) and methyl cyanide (_Codella et al._ , 2009) (right).
The colour images refer to formamide for the left and middle panels and to
methyl cyanide for the rightmost one. Red circles show the colour images
synthetic beams. The beams of the HDCO, H2CO, and CH3CHO observations are
equal to that of formamide

. Time increases and chemistry evolves going from North to South. The youngest
shocked region (within L1157-B1) is that imaged by the SiO emission at high
velocity.

The new generation of broad-band heterodyne receivers has revolutionized the
field by significantly reducing the observation time cost of unbiased spectral
line surveys in the (sub)millimeter domain making it possible new molecule
detections towards L1157-B1: acetaldehyde and its D-isotopomer, CH2DOH,
formamide (_Sugimura et al._ , 2011; _Yamaguchi et al._ , 2012), dimethyl
ether, ketene (H2CCO) and glycolaldehyde (_Lefloch et al._ , 2018), within the
IRAM Large Program ASAI (see § 3.3: _Lefloch et al._ , 2017). Remarkably, the
same ASAI survey towards the source driving the L1157 outflow, L1557-mm,
showed that the latter possesses a much poorer spectrum, dominated by
C-bearing species (see § 4.2) in contrast with the iCOM rich spectrum at the
shocked gas site B1 (_Bachiller et al._ , 2001).

The list of iCOMs detected in shocks is summarized in Table 3. The ASAI
IRAM-30m spectral survey provides beam-averaged ($\sim$ 30$\arcsec$) column
densities towards L1157-B1 (_Lefloch et al._ , 2017): $\sim$ 1015 cm-2
(CH3OH), 2–5 $\times$ 1013 cm-2 (CH3COH, HCOOCH3, HCOCH2OH, CH3OCH3, C2H5OH),
and $\sim$ 1012 cm-2 (NH2CHO). Thanks to a multiline CO analysis, _Lefloch et
al._ (2012) estimated in the same region an H2 column density equal to $\sim$
1021 cm-2. In turn, this implies iCOMs abundances in the range from a few 10-9
to $\sim$ 10-7 (CH3OH). The spatial distributions obtained thanks to IRAM
NOEMA (_Codella et al._ , 2017, 2020) reveals higher values at the peak
regions (not beam-diluted): $\sim$ 1016 cm-2 (CH3OH), 7 $\times$ 1013 cm-2
(CH3COH), $\sim$ 1013 cm-2 (NH2CHO). In the same vein, imaging the
NGC1333-IRAS4A shocked regions, (_De Simone et al._ , 2020b) derived up to 5
$\times$ 1015 cm-2 for methanol, $\sim$ 1014 cm-2 for acethaldeyde and
dimethyl ether, and 5 $\times$ 1012 cm-2 for formamide.

When considering the case of L1157-B1, where the largest number of iCOMs is
detected, there is a predominance of O-bearing iCOMs, seven detected, against
only three and one N- and S- bearing iCOMs, respectively (see Tab. 3). In
general, the abundances of iCOMs other than methanol are similar, $\sim
10^{-8}$, within a factor of 3, and represent about 2–5% of the methanol
abundance. Puzzlingly, relatively large iCOMs such as acetone, ethylene
glycol, acrylonitrile and ethyl cyanide have not been detected so far, despite
they are commonly observed towards hot corinos. What makes this lack of
detection more puzzling is that, on average, iCOM abundances in L1157-B1 are
larger than those measured in hot corinos by a factor 2-10, depending on the
species, with the possible exception of dimethyl ether (_Lefloch et al._ ,
2017). _Arce et al._ (2008) were the first to notice that the iCOM abundances
relative to methanol are similar to that found in massive hot cores and
molecular clouds in the Galactic Center region. Assuming that iCOMs form on
the grain surfaces, these authors suggested that the dust mantle composition
in the L1157 dark cloud could be similar to that in hot cores and the Galactic
Center molecular clouds.

Finally, _Holdship et al._ (2019b) analysed the line profiles of a sample of
eight molecular outflow shocks. They found an approximately constant abundance
ratio of methanol over acetaldehyde in the lower-velocity gas and suggested
that this is in favor of a grain surface scenario for acetaldehyde formation.
However, we note that, if acetaldehyde is formed by the ethanol (_Vazart et
al._ , 2020) which is probably a grain-surface product (_Perrero et al._ ,
2022), a constant abundance ratio [CH3OH]/[CH3CHO] would likely result too.
_Holdship et al._ (2019b) also found evidence for the destruction of methanol
and acetaldehyde in the post-shock gas of fast outflows (see also _Suutarinen
et al._ , 2014). We caution, though, that these results need to be verified
using high-spatial-resolution maps of shocked regions.

###  6.2 The post-shock gas chemical structure

It has recently become possible with radio to millimeter interferometers to
resolve the molecular gas emission down to spatial scales of a few 100 au or
less in the nearest star forming regions. This instrumental progress has
allowed to peer into the time-dependent chemical structure of the shocked gas
and to set stringent constraints on the possible iCOM formation pathway(s) at
work in shocked regions.

As already mentioned, the best characterized source is the L1157-B1 shock
region, in which several iCOMs and closely related species were mapped at
sufficiently high angular resolution to constrain their possible formation
pathways (Fig. 13 and Table 3). Dynamical modelling of the outflow and its
precessing jet (_Gueth et al._ , 1996; _Podio et al._ , 2016) has provided a
detailed picture of the mass-ejection history and the associated shock
regions, which trace the jet impact against the outflow cavity. Based on these
studies, the kinematical age of the L1157-B1 region is $\sim$1000-2000 yr.
Figure 13 shows the chemical differentiation observed in L1157-B1 using the
NOEMA interferometer, particularly between formamide and acetaldehyde. A key
to interpret the L1157-B1 observed chemical structure is SiO, expected to form
in high-velocity ($\geq$ 10 km s-1) shocks able to sputtering dust mantles as