11institutetext: European Southern Observatory, Karl-Schwarzschild-Strasse 2,
85748 Garching bei München, Germany
11email<EMAIL_ADDRESS>22institutetext: Leiden Observatory, Leiden
University, PO Box 9513, 2300RA, Leiden, The Netherlands 33institutetext: Max-
Planck-Institut für Extraterrestrische Physik, Giessenbachstrasse 1, D-85748
Garching, Germany 44institutetext: Department of Astronomy, University of
Michigan, 500 Church Street, Ann Arbor, MI 48109, USA 55institutetext:
Institute of Astronomy and Astrophysics, Academia Sinica, No. 1, Sec. 4,
Roosevelt Road, Taipei 10617, Taiwan, R. O. C. 66institutetext: National
Astronomical Observatory of Japan, NAOJ Chile, Alonso de Córdova 3788, Office
61B, 7630422, Vitacura, Santiago, Chile 77institutetext: Joint ALMA
Observatory, Alonso de Córdova 3107, Vitacura, Santiago, Chile
88institutetext: Anton Pannekoek Institute for Astronomy, University of
Amsterdam, Science Park 904, 1098 XH Amsterdam, The Netherlands
99institutetext: Star and Planet Formation Laboratory, RIKEN Cluster for
Pioneering Research, Wako, Saitama 351-0198, Japan 1010institutetext: National
Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, VA 22903, USA

# Which molecule traces what:
chemical diagnostics of protostellar sources

Łukasz Tychoniec  Ewine F. van Dishoeck 1122 Merel L.R. van ’t Hoff 2233
Martijn L. van Gelder  4 4 Benoît Tabone 22
Yuan Chen 22 Daniel Harsono 22 Charles L. H. Hull 55 Michiel R. Hogerheijde
6677 Nadia M. Murillo 22 8 8 John J. Tobin 991010

###### Abstract

Context.  The physical and chemical conditions in Class 0/I protostars are
fundamental in unlocking the protostellar accretion process and its impact on
planet formation.

Aims. The aim is to determine which physical components are traced by
different molecules at sub-arcsecond scales (<100 – 400 au).

Methods. We use a suite of Atacama Large Millimeter/submillimeter Array (ALMA)
datasets in Band 6 (1 mm), Band 5 (1.8 mm) and Band 3 (3 mm) at spatial
resolutions 0$\aas@@fstack{\prime\prime}$5 – 3″ for 16 protostellar sources.
For a subset of sources, Atacama Compact Array (ACA) data at Band 6 with a
spatial resolution of 6′′ are added. The availability of both low- and high-
excitation lines, as well as data on small and larger scales, are important to
unravel the full picture.

Results. The protostellar envelope is well traced by C18O, DCO+ and N2D+, with
the freeze-out of CO governing the chemistry at envelope scales. Molecular
outflows are seen in classical shock tracers like SiO and SO, but ice-mantle
products such as CH3OH and HNCO released with the shock are also observed. The
molecular jet is a key component of the system, only present at the very early
stages, and prominent not only in SiO and SO but also occasionally in H2CO.
The cavity walls show tracers of UV-irradiation such as hydrocarbons C2H and
c-C3H2 as well as CN. The hot inner envelope, apart from showing emission from
complex organic molecules (COMs), also presents compact emission from small
molecules like H2S, SO, OCS and H13CN, most likely related to ice sublimation
and high-temperature chemistry.

Conclusions. Sub-arcsecond millimeter-wave observations allow to identify
those (simple) molecules that best trace each of the physical components of a
protostellar system. COMs are found both in the hot inner envelope (high
excitation lines) and in the outflows (lower-excitation lines) with comparable
abundances. COMs can coexist with hydrocarbons in the same protostellar
sources, but they trace different components. In the near future, mid-IR
observations with JWST–MIRI will provide complementary information about the
hottest gas and the ice mantle content, at unprecedented sensitivity and at
resolutions comparable to ALMA for the same sources.

###### Key Words.:

stars: formation, astrochemistry, techniques: interferometric, ISM: molecules,
submillimeter: ISM

##  1 Introduction

The formation of Sun-like stars is set in motion by the collapse of a cold,
dense cloud. Many different physical processes take place in the protostellar
stage – first few $10^{5}$ yrs that are critical to the subsequent evolution
of the star and its planetary system (Lada 1987). The mass of the star and
that of its circumstellar disk are determined during this embedded phase
(Hueso & Guillot 2005) and the first steps of planet formation must take place
then (Greaves & Rice 2010; Williams 2012; ALMA Partnership et al. 2015; Manara
et al. 2018; Harsono et al. 2018; Tobin et al. 2020; Tychoniec et al. 2018,
2020; Segura-Cox et al. 2020). At the same time, on larger scales, the
collapsing envelope is dispersed by the energetic action of bipolar jets and
winds emanating from the star-disk system which create outflows of entrained
gas and dust (Arce & Sargent 2006; Offner & Arce 2014)

Rotational transitions of molecules are a powerful tool to probe other
components of the system and can be used to infer densities, temperatures, UV
fields, chemical abundances and kinematics (van Dishoeck & Blake 1998; Evans
1999). Until recently, studies of low-mass protostars have suffered from
insufficient spatial resolution to disentangle these different physical
components. The advent of submillimeter interferometry opened the possibility
to study protostellar systems at much smaller scales than with single-dish
observations (e.g., Chandler & Carlstrom 1996; Hogerheijde et al. 1999;
Schilke et al. 1992; Wilner et al. 2000; Jørgensen et al. 2005a; Tobin et al.
2011).

With the Atacama Large Millimeter/submillimeter Array (ALMA), it is possible
to image many molecular lines on the relevant physical scales with achievable
observing times at sub-arcsecond resolution. Impressive ALMA studies of
individual low-mass protostars have been presented, focusing both on simple
species (¡ 6 atoms) and complex molecules (¿ 6 atoms) (e.g., Sakai et al.
2014a; Jørgensen et al. 2016; López-Sepulcre et al. 2017; Lee et al. 2019a;
Codella et al. 2018; Lee et al. 2019b; Manigand et al. 2020; van Gelder et al.
2020; Bianchi et al. 2020). The first ALMA surveys of complex molecules on a
larger sample are taking place (Yang et al. 2021) and of Class I disks (Podio
et al. 2020; Garufi et al. 2021). The CALYPSO survey with NOEMA studied a
larger sample of protostars but at more limited resolution and sensitivity
(Belloche et al. 2020; Maret et al. 2020; Podio et al. 2021).

Here we present ALMA data of 16 protostellar sources covering rotational
transitions of various molecules; we use these data to build a complete
picture of what types of molecules trace which physical structures in
protostars. This sample constitutes one of the largest combinations of high-
resolution ALMA observations of Class 0/I protostars to date in ALMA Band 3,
5, and 6. Covering a broad range of protostellar properties within the low-
mass regime, the aim is to identify and describe key molecular tracers of
future Sun-like stars and what physical components of star-forming sources
they correspond to. Parts of the ALMA datasets presented here have been
already published, focusing on different aspects: complex organic molecules
(van Gelder et al. 2020; Nazari et al. 2021); Class I disks temperature
structure (van ’t Hoff et al. 2020a); outflows and high-velocity jets in
Serpens (Hull et al. 2016; Tychoniec et al. 2019), molecular emission
associated with magnetic fields (Hull et al. 2017; Le Gouellec et al. 2019).
In this work we make a comprehensive overview of these different datasets,
making full use of those observations, with uniform analysis methods. This
allows to reveal and systematise the molecular tracers of Class 0/I
protostars. Out of 16 presented sources 11 are included in upcoming James Webb
Space Telescope (JWST) observations with the Mid-Infrared Instrument (MIRI;
Wright et al. 2015) (GTO1290; GTO1257) and Near-Infrared Spectrograph
(NIRSpec) ( GO1960; GTO1186; GO1798; GO2104)

###  1.1 Physical components of a protostellar system

Class 0 sources are defined by their strong excess of submillimeter luminosity
and very low bolometric temperatures <70 K (André et al. 1993; Chen et al.
1995). These sources are associated with powerful outflows, and the envelope
mass dominates the mass of the entire system. Class I sources are defined by
having an infrared spectral index indicating strong reddening (Lada 1987) with
bolometric temperatures of 70–650 K (Chen et al. 1995). Those systems have
already converted most of their envelope mass into disk and protostar (Crapsi
et al. 2008; van Kempen et al. 2009b; Maury et al. 2011). For the typical
envelope masses of sources presented here and average disk masses found by
Tychoniec et al. (2020), the $M_{\rm disk}/M_{\rm env}\simeq$ 1% for Class 0
and $\simeq$ 20% for Class I, with values up to 75–98% in cases of
rotationally supported disks (Jørgensen et al. 2009).

Table 1: Targeted protostellar systems

Source name | R.A. | Decl. | $d$ | Class | $L_{\rm bol}$ | $T_{\rm bol}$ | $M_{\rm env}$ | Ref.
---|---|---|---|---|---|---|---|---
| (J2000) | (J2000) | (pc) |  | (L⊙) | (K) | (M⊙) |
Serpens SMM1 | 18:29:49.8 | +01:15:20.5 | 439 | 0 | 109 | 39 | 58 | (1)
Serpens S68N | 18:29:48.1 | +01:16:43.3 | 439 | 0 | 6 | 58 | 10 | (2)
Ser-emb 8 (N) | 18:29:48.7 | +01:16:55.5 | 439 | 0 | — | — | — | —
Serpens SMM3 | 18:29:59.2 | +01:14:00.3 | 439 | 0 | 28 | 38 | 13 | (1)
BHR 71 | 12:01:36.3 | -65:08:53.0 | 200 | 0 | 15 | 44 | 2.7 | (1)
IRAS 4B | 03:29:12.0 | +31 13 08.1 | 293 | 0 | 7 | 28 | 4.7 | (1)
Per-emb-25 | 03:26:37.5 | +30:15:27.8 | 293 | 0/I | 1.9 | 61 | 2.0 | (2)
B1-c | 03:33:17.9 | +31:09:31.8 | 293 | 0 | 5 | 48 | 15 | (2)
HH211-mm | 03:43:56.8 | +32:00:50.2 | 293 | 0 | 2.8 | 27 | 19 | (2)
L1448-mm | 03:25:38.9 | +30:44:05.3 | 293 | 0 | 13 | 47 | 15 | (2)
L1527 IRS | 04:39:53.9 | +26:03:09.5 | 140 | 0/I | 1.6 | 79 | 0.12 | (3)
B5-IRS1 | 03:47:41.6 | +32:51:43.7 | 293 | I | 7 | 181 | 3.5 | (2)
TMC1 | 04:41:12.7 | +25:46:34.8 | 140 | I | 0.9 | 101 | 0.14 | (1)
IRAS 04302 | 04:33:16.5 | +22:53:20.4 | 140 | I | 0.7 | 300 | 0.05 | (1)
L1489 IRS | 04:04:43.0 | +26:18:57.0 | 140 | I | 3.8 | 200 | 0.2 | (1)
TMC1A | 04:39:34.9 | +25:41:45.0 | 140 | I | 2.7 | 118 | 0.2 | (1)

  * •

(1) Kristensen et al. 2012, (2) Enoch et al. 2009, (3) Green et al. 2013

The different components of protostellar systems vary significantly in their
physical conditions, such as density and temperature, molecular enrichment,
and dynamics. Our current knowledge about them is described briefly below to
set the scene for the interpretation of our data.

Envelope. The envelope surrounding a protostar is the material that fuels the
accretion process onto the star and disk. The physical conditions in the outer
envelope on scales of a few 1000 au are reminiscent of those of starless cores
with heavy freeze-out, and their chemical composition is directly inherited
from the cloud out of which the star is being born (Caselli & Ceccarelli
2012). Systematic motions such as infall or expansion can occur but otherwise
they are characterized by low turbulence and narrow (FWHM ¡ 0.5 - 1 km s-1)
line profiles indicative of quiescent gas (Jørgensen et al. 2002).

Warm inner envelope. In the innermost part of the envelope on scales of the
disk, temperatures rise above 100 K, so any water and complex organic
molecules (COMs) contained in ices are released from the grains back into the
gas where they are readily observed at submillimeter wavelengths. This region
with its unique chemical richness is called the hot core, or to distinguish it
from its high-mass counterpart, hot corino (Herbst & van Dishoeck 2009).

Jets and outflows. As the material is accreting from the envelope onto the
disk, excess angular momentum has to be transported via a still elusive
process to allow material to accrete on the growing protostar. Jets and
outflows constitute compelling candidates to extract angular momentum via
magnetic fields. In the earliest stages when the mass loss is at its peak, the
densities are high enough to form molecules in the internal shocks in the jet
(Bachiller & Gomez-Gonzalez 1992; Tafalla et al. 2010). Much slower (<20 km
s-1) and less collimated gas moving away from the protostar is called an
outflow. Their origin remains debated. Large scale outflows reveal bow-shock
shells and cavities possibly driven by the fast intermittent jet (Gueth et al.
1996; Gueth & Guilloteau 1999; Tychoniec et al. 2019). Temperatures in shocked
regions are much higher than in the surrounding envelope, up to a few thousand
K, and sputtering of grain cores and ice mantles can further result in unique
chemical signatures (Arce et al. 2008; Flower & Pineau des Forêts 2013).

Outflow cavity walls. These are the narrow zones in between the cold dense
quiescent envelope material and the lower-density warm cone where outflows are
propagating at large velocities. Cavity walls are exposed to UV radiation from
the accreting star-disk boundary layer, which can escape through the outflow
cavity without being extincted (Spaans et al. 1995). This creates conditions
similar to those found in Photon Dominated Regions (PDRs), which occur
throughout the interstellar medium near sources of intense UV radiation
(Hollenbach & Tielens 1997). In units of the interstellar radiation field
(ISRF, Draine 1978), typical values of 102–103 are found on scales of $\sim
1000$ au (van Kempen et al. 2009a; Yıldız et al. 2012; Benz et al. 2016;
Karska et al. 2018).

Young disk. In the inner envelope, a protoplanetary disk starts to form as the
natural outcome of a rotating collapsing core (Ulrich 1976; Cassen & Moosman
1981; Terebey et al. 1984). A young disk should be rotating in Keplerian
motion. At early stages it is difficult to identify whether the so-called
embedded disk is rotationally supported, since any molecular emission from the
disk is entangled with that from the envelope. In recent years several
embedded disks have been identified to have Keplerian rotational structure on
scales of $\sim$100 au (Tobin et al. 2012; Murillo et al. 2013; Ohashi et al.
2014; Yen et al. 2017). Molecular tracers in young disks, apart from providing
the kinematic information, can probe their temperature structure as well (van
’t Hoff et al. 2018b).

This work is organized as follows. In Section 2, the observations used in this
work are presented, Section 3 presents the results of this work including
detections and morphology of the targeted molecules. In Section 4 the results
are discussed. with special focus on which molecular tracers are corresponding
to each of the physical components. The focus is on a qualitative description,
rather than quantitative analyses for which source specific models and more
rotational transitions of a given molecule would be needed. We summarize our
work in Section 5.

##  2 Observations

###  2.1 Datasets

Six different ALMA 12m datasets at Band 3, 5, and 6 are used in this work to
cover 14 out of 16 sources. The spatial resolution of all ALMA 12m datasets is
comparable (0$\aas@@fstack{\prime\prime}$3–0$\aas@@fstack{\prime\prime}$6),
except for 2017.1.01174.S, where Band 3 observations are obtained at 3″.
Additionally for 6 out of 16 sources in Band 6, ACA observations with 7m
antennas were obtained at 6″ resolution.

Figure 1:  Continuum emission at 1.3 mm of four example protostellar systems
obtained with ALMA at 0$\aas@@fstack{\prime\prime}$5 resolution. Symbols of
stars point to confirmed protostellar sources, circles show condensations of
continuum emission, without confirmed protostellar nature, dotted lines show
outflow cavity walls, and dashed lines show streams of envelope material.
Arrows indicate outflow directions.

The targeted protostars are well-known objects located in different star-
forming regions. They span a range of properties within the low-mass regime;
the probed range of $L_{\rm bol}$, $T_{\rm bol}$ and $M_{\rm env}$ is shown in
Table 1, and are provided by a suite of observations across the infrared and
submillimeter spectrum (Enoch et al. 2009; Kristensen et al. 2012; Green et
al. 2013).

The details of the observations are summarized in Table 4. The spatial
resolution allows to observe protostellar systems at solar-system scales; Band
5 and 6 observations provide a resolution of
$\sim$0$\aas@@fstack{\prime\prime}$5, which corresponds to a 70–220 au
diameter for sources in our sample. Thus, regions down to 35-110 au radius in
the inner envelope are probed. The Band 3 data achieve moderate resolution of
$\sim$3″ which provides information on intermediate envelope scales of
500–1500 au. The ACA observations of six sources at 6″ resolution probe
envelope scales of 800–2000 au.

###  2.2 Spectral setup of the observations

A collection of different datasets using different ALMA bands implies varying
spectral and spatial resolution as well as spectral coverage across the
analysis. This is the reason that throughout this paper the sources shown in
the figures differ when presenting detections and maps of different molecules.
In all cases, when the molecule is discussed, only those sources where the
given transition has been targeted are discussed. All non-detections are
explicitly stated. Table LABEL:table:table_mols provides a list of targeted
molecular transitions, with sources that have a particular line covered and
detected or not detected.

Figure 2: Maps of key envelope tracers toward B1-c (Class 0, top) and TMC1
(Class I, bottom) obtained with ACA. Contours represent continuum emission at
1.3 mm observed with ACA. Note different distances to B1-c and TMC1 resulting
in different spatial resolutions of the maps. Left: C18O 2 – 1 Middle: N2D+ 3
– 2. Right: DCO+ 3 – 2. All moment 0 maps are integrated from -2.5 to 2.5 km
s-1 w.r.t $\varv_{\rm sys}$.

The ALMA observations presented here target different spectral setups across
Band 3, 5, and 6. This allows to probe the strongly varying physical scales
and conditions. In particular, our Band 3 data grants access to lines at very
low excitation levels that enable tracing more extended material. Dust is less
optically thick in Band 3 compared with Band 6 which potentially allows to
peek inside the densest inner regions. Band 6 offers a usual set of tracers of
outflows: CO, SiO and SO.

Cold outer envelopes with temperatures < 20 K are probed with low $E_{\rm up}$
transitions. Additionally, non-thermal processes such as sputtering of
material from the grains in the outflow, will also be seen in low $E_{\rm up}$
due to their lower critical densities. On the other hand, thermal desorption
from grains in the innermost regions are best probed with lines with high
$E_{\rm up}$. With the large span of frequencies of the observations,
different transitions of the same molecule can be detected and used to trace
different components of the system (e.g., a HNCO line at $E_{\rm up}=$ 15 K is
available in Band 3 and lines at 70 and 125 K are covered in Band 6).

##  3 Results

###  3.1 Continuum emission from protostars

Figure 1 presents continuum emission maps toward four example protostars
obtained with ALMA at $\sim$ 0$\aas@@fstack{\prime\prime}$5\. The examples
illustrate most characteristic features observed in the continuum maps. The
continuum images for all sources are presented in Fig.
LABEL:fig:continuum_no1_v1.

The continuum emission observed at millimeter wavelengths (1.3 - 3 mm in our
observations) traces thermal dust emission from the inner envelope and the
embedded disk. In the Class 0 protostars, the central continuum source
generally appears compact $R$ <100 au (e.g., B1-c, Fig. 1). The example of
SMM3 (Fig. LABEL:fig:continuum_no1_v1) shows a large resolved dust structure
perpendicular to the outflow, but its classification as a disk is not certain.
The fact that we observe primarily compact continuum emission towards Class 0
sources is consistent with observations of confirmed rotationally-supported
disks (e.g., Tobin et al. 2018; Maury et al. 2019; Tobin et al. 2020) and
predictions of models (Visser & Dullemond 2010; Harsono et al. 2015b; Machida
et al. 2016). While we assume that this compact emission belongs to the young,
embedded disk, we do not have spatial and kinematic resolution to confirm the
presence of Class 0 Keplerian disks in our sample.

The extended continuum emission in Class 0 sources is consistent with a
significant amount of envelope material surrounding the protostar. In the case
of the SMM1 system, presented in Fig. 1 (left), the continuum clumps outside
the central emission are components of multiple protostellar systems, marked
with stars, confirmed by the presence of individual molecular outflows (Hull
et al. 2016, 2017). Binary components are also seen in TMC1, BHR71, and IRAS
4B (Fig. LABEL:fig:continuum_no1_v1). In the case of S68N presented in Fig. 1,
two emission peaks that stand out from the diffuse envelope emission are
marked with circles, but their protostellar nature is not confirmed.

###  3.2 Protostellar envelope

We present molecules that trace the bulk of the protostellar envelope in Fig.
2. The protostellar envelope has a typical radius on the order of a few 1000
au (Jørgensen et al. 2002; Kristensen et al. 2012). Thus, the sub-arcsecond
ALMA 12m array observations tend to resolve-out the envelope emission. For
instance, the maximum recoverable scale (MRS) of ALMA Band 6 observations at
0$\aas@@fstack{\prime\prime}$4 presented here is 5″, which is between 600–2000
au diameter depending on distance to the source. For that reason, we discuss
in this section mainly the ALMA-ACA observations obtained at lower spatial
resolution (6″; 750–2500 au) for six sources in our datasets; Class 0 sources:
B1-c, BHR71, Per-emb-25, SMM3, and IRAS 4B, and Class I source TMC1. The ACA
can zoom-in on what was previously contained in a single-dish beam of 15–20″,
while the MRS of ACA (30″) enables us to preserve sensitivity to large-scale
emission. The MRS of all observations presented here are reported in Table 4.

In Fig. 2 typical envelope tracers C18O 2–1 ($E_{\rm up}=16$ K), DCO+ 3–2
($E_{\rm up}=21$ K), and N2D+ 3–2 ($E_{\rm up}=22$ K), observed at 6″
resolution with the ACA, are presented toward example Class 0 and Class I
sources – B1-c and TMC1, respectively. The emission from the presented
molecules exhibits similar behaviour for all Class 0 sources, therefore B1-c
serves as a representative case; TMC1 is the only Class I source in the sample
with 7m observations available. The maps for all sources for which these
molecules have been targeted can be found in Appendix LABEL:AppendixD. All
envelope tracers presented here are characterized by narrow line profiles with
FWHM $\sim 1$ km s-1.

The C18O emission peak coincides with the continuum peak for our six sources
and appears to be compact, less than 1000 au diameter for B1-c and TMC1. For
B1-c and all Class 0 sources (Fig. LABEL:fig:envelope_co18_7m), low-level
extended C18O emission is seen along the outflow direction. For the only Class
I source targeted with the ACA, emission is marginally resolved in the
direction perpendicular to the outflow.

In our observations N2D+ is seen extended in the direction perpendicular to
the core major axis toward B1-c (Fig. 2) and other Class 0 sources except IRAS
4B (Fig.LABEL:fig:envelope_n2d+_7m). In the case of IRAS 4B the emission from
this molecule appears dominated by large-scale emission from the filament
detected toward this source, connecting it with IRAS 4A (Sakai et al. 2012).
The peak of the N2D+ emission is significantly shifted from the continuum peak
in all cases, with a significant decrease in the inner regions in some cases
(see BHR 71 in Fig. LABEL:fig:envelope_n2d+_7m). Similar extended N2D+
emission in other Class 0 sources was seen by Tobin et al. (2013) based on
lower resolution SMA and IRAM-30m data. For TMC1 the N2D+ molecule is not
detected.

The DCO+ emission is seen extended in a similar fashion to what is observed
for N2D+. However, contrary to N2D+, DCO+ is brightest on the continuum peak
for all sources except TMC1 and Per-emb-25. For these two sources, the
emission peak is offset by 1000–2000 au from the continuum source in the
direction perpendicular to the outflow. In the Class I source TMC1, DCO+ is
present on much smaller scales ($<2000$ au radius) than in Class 0 sources.

Figure 3: Maps of the EHV jets observed in SiO. Moment 0 maps are presented in
color scale with continuum emission at 1.3 mm presented in black contours,
both obtained with 12m observations. SiO 4–3 map of HH211 integrated from -20
to -10 and from 10 to 20 km s-1 w.r.t $\varv_{\rm sys}$; B1-c ntegrated from
-70 to -40 and from 40 to 70 km s-1 w.r.t $\varv_{\rm sys}$; SMM3 integrated
from -60 to -40 and from 20 to 35 km s-1 w.r.t $\varv_{\rm sys}$.  Figure 4:
Zoom-in on molecular bullets from the SMM3 jet (see Fig. 3). CO, SiO, SO, and
H2CO molecular transitions are presented. Top: Norther/blueshifted bullet.
Moment 0 maps are integrated from -60 to -40 km s-1 w.r.t $\varv_{\rm sys}$.
The map center is offset from the SMM3 continuum center by
(–3$\aas@@fstack{\prime\prime}$7, +10$\aas@@fstack{\prime\prime}$3). Bottom:
Southern/redshifted bullet. Moment 0 maps are integrated from 20 to 40 km s-1.
The map center is offset from the SMM3 continuum center by
(+2$\aas@@fstack{\prime\prime}$7, –7$\aas@@fstack{\prime\prime}$). Figure 5:
Low-velocity outflow in CO, SO, and SiO. Moment 0 maps of CO toward TMC1, SO
and SiO map toward B1c are integrated from -10 to 10 km s-1 w.r.t $\varv_{\rm
sys}$.

###  3.3 Outflows and jets

Figure 3 presents the extremely high-velocity (EHV) molecular jet component
for HH211, B1c, and SMM3 observed in SiO with the ALMA 12m array; the L1448-mm
EHV is shown in Fig. 12. The latter and HH211 are well-known EHV sources
(Guilloteau et al. 1992; Lee et al. 2007) while SMM3 and B1-c are new
detections of the jet component. HH211 shows SiO emission at low-velocities
because the outflow is almost in the plane of the sky, but the high velocities
are evident from the large proper motion movements of the bullets ($\sim$115
km s-1; Lee et al. 2015). CO, H2CO, and SiO in the EHV jets of Emb8N and SMM1
are presented in detail in Hull et al. (2016) and Tychoniec et al. (2019).

These data are particularly interesting as still only few molecules tracing
EHV jets have been identified to date (see Lee 2020 for review). Apart from
those presented here – CO, SiO, SO, and H2CO– molecules such as HCO+ and H2O
have been seen in this high-velocity component (Kristensen et al. 2012; Lee et
al. 2014). What is especially important to highlight is the third detection of
a H2CO bullet in SMM3 after IRAS 04166 (Tafalla et al. 2010) and Emb8N
(Tychoniec et al. 2019). These detections mean that either a significant
fraction of ice-coated dust is released with the jet, or that the H2CO is
efficiently produced in the jet through gas-phase chemistry.

In the case of L1448-mm and HH211, there are several molecular bullets along
the jet axis with velocities up to 100 km s-1, while both B1-c and SMM3 show a
much simpler structure with two bullets detected on one side in the former and
a single pair of symmetrically placed bullets observed in the latter case.
B1-c actually has a pair of bullets $\sim$ 200 au from the continuum peak,
with the other bullet at 2500 au only seen in the redshifted part of the jet.
The emission in SiO and SO appears very similar (Fig.
LABEL:fig:sio_outflows_no2, Fig. LABEL:fig:so_outflow). There are no other
molecules tracing the high-velocity component toward this source. H2CO and
12CO are not targeted with our ALMA 12m datasets toward B1-c. The SMM3 jet has
two distinct high-velocity bullets at $\sim$ 3200 au from the source which
appear similar in CO, SiO and SO (see zoom-in on Fig 4). Additionally, the
redshifted bullet shows faint, but significant emission from H2CO. No traces
of H2CO are found in the blueshifted outflow.

Figure 5 presents low-velocity outflow tracers CO 2–1 ($E_{\rm up}=16$ K) for
TMC1, SO 56–45 ($E_{\rm up}=35$ K), and SiO 4–3 ($E_{\rm up}=21$ K) for B1-c.
In Fig. LABEL:fig:co_outflows we present an overview of CO 2–1 emission for
five sources obtained with the ALMA 12 m array at
0$\aas@@fstack{\prime\prime}$4 resolution. The sources show a variety of
emission structures in the low velocity gas (<20 km s-1). In all cases, we do
not capture the entirety of the outflows as they extend beyond the primary
beam of observations ($\sim$ 30″). SMM3 and Emb8N have very narrow outflow
opening angles ($<$ 20 degrees), while SMM1, S68N and TMC1 present larger
opening angles.

CO emission is especially prominent in the cavity walls, which can be related
to both the limb brightening effect as well as higher (column) density of the
material in the outflow cavity walls. This is especially highlighted in the
Class I source, TMC1, where CO emission is almost exclusively seen in the
outflow cavity walls. Even though CO is piling up in this region, it is
observed at velocities up to 15 km s-1 so it is clearly tracing the entrained
material and not the envelope. The lower envelope density in Class I results
in less material to be entrained in the outflow. In SMM1, three CO outflows
from SMM1-a, SMM1-b, and SMM1-d are overlapping (Hull et al. 2016; Tychoniec
et al. 2019).

In Fig. LABEL:fig:sio_outflows_no2, SiO maps are presented: Band 5 SiO 4–3
($E_{\rm up}=21$ K) and Band 6 SiO 5–4 ($E_{\rm up}=31$ K) data are shown in
velocity ranges corresponding to the low-velocity outflow. In contrast to CO
emission, the low-velocity SiO is mostly observed in clumps of emission
instead of tracing the entirety of the outflowing gas. Several such clumps can
be seen in the S68N source. In some cases, the clumps are relatively symmetric
(Emb8N, B1-c), while monopolar emission is seen in other examples (L1448-mm,
SMM1-d). In the case of SMM1-a and SMM1-b, very weak SiO emission at low
velocities is observed. SiO emission in outflows is exclusively present in the
Class 0 sources, while absent in the Class I sources, TMC1 and B5-IRS1,
covered in these data sets.

Fig. LABEL:fig:so_outflow presents SO 56–45 ($E_{\rm up}=35$ K) and SO 67–56
($E_{\rm up}=47$ K) observations in Band 6. The emitting regions of SO are
comparable with those of SiO for the Class 0 sources. The cases of S68N and
B1-c show that SO emission also peaks at the source position while SiO is
absent there. Thus, SO and SiO do not always follow each other and some SO
might be associated with hot core emission (Drozdovskaya et al. 2018).
Important differences are observed for TMC1, where SO seems to be associated
with the remainder of the envelope or the disk, while the SiO is not detected
toward this source, as mentioned above.

HCN 1–0 ($E_{\rm up}=4$ K) and H13CN 2–1 ($E_{\rm up}=12$ K) maps are
presented in Fig. LABEL:fig:hcn_outflow and LABEL:fig:h13cn, respectively. HCN
is clearly seen in outflowing material enhanced in similar regions as low-
velocity SiO. For Emb8N the HCN emission has been associated with intermediate
velocity shock (Tychoniec et al. 2019). In B1-c and L1448-mm weak extended
emission along the outflow direction is detected but H13CN strongly peaks on
source. In the case of HH211, H13CN is seen only in the outflow, with a
geometry consistent with the outflow cavity walls, but with velocity profiles
that are consistent with the outflowing material.

Figure 6: Maps of the ice mantle tracers toward the S68N outflow, with the CO
low-velocity outflow map for reference. Top: CO, H2CO and CH3OH moment 0 maps
obtained in Band 6 at 0$\aas@@fstack{\prime\prime}$5 resolution. Circles show
regions from which spectra were obtained for analysis in Section 7.1. Bottom:
CH3CHO, HNCO, and CH3CN moment 0 maps obtained in Band 3 at
2$\aas@@fstack{\prime\prime}$5 resolution. The emission is integrated from -10
to -1 km s-1 and from 1 to 10 km s-1 w.r.t $\varv_{\rm sys}$. Figure 7: Maps
of the outflow cavity wall tracers toward SMM3 and B1c, with low-velocity
outflow map for reference. Top: Moment 0 maps toward SMM3 of CO 2–1, c-C3H2
61,6–50,5 , and 13CS 5–4 obtained in Band 6 at 0$\aas@@fstack{\prime\prime}$5
resolution and CN 1–0 in Band 3 at 3″with continuum emission at the same band
and resolution in black contours. Bottom: Moment 0 maps toward B1c of SO
67–56, C2H 32.5,3–21.5,1 and H13CO+ 3–2 obtained at
0$\aas@@fstack{\prime\prime}$5 and CN 1–0 at 3″with continuum emission at the
same band and resolution in black contours. The emission is integrated from -5
to -1 km s-1 and from 1 to 5 km s-1 w.r.t $\varv_{lsr}$. Outflow directions
and delineated cavity walls are shown in C2H and C3H2 maps.

Ice-mantle tracers are a different class of molecules detected in low-velocity
protostellar outflows. They are produced and entrained through interactions
between the jet and the envelope. Here we present ALMA 12m array observations
in Band 6 at 0$\aas@@fstack{\prime\prime}$5 resolution for CO 2 – 1 ($E_{\rm
up}=17$ K), CH3OH 21,0 – 10,0 ($E_{\rm up}=28$ K), and H2CO 30,3 – 20,2
($E_{\rm up}=21$ K) , and in Band 3 at 3″ for CH3CN 61 – 51 ($E_{\rm up}=26$
K), CH3CHO 61,6,0 – 51,5,0 ($E_{\rm up}=21$ K), and HNCO 50,5 – 50,4 ($E_{\rm
up}=16$ K). Fig. 6 compares maps of integrated emission from those molecules
with those of CO for S68N. All ice-mantle tracers detected in the outflow are
observed in their low-energy transitions. Additional maps for S68N are
presented in the Appendix (Fig. LABEL:fig:fig9_extra): Band 6
0$\aas@@fstack{\prime\prime}$5 resolution image of CH3CHO 140,14–130,13
($E_{\rm up}=96$ K) and H2CCO 131,13–121,12 ($E_{\rm up}=101$ K), which is
overlapping with NH2CHO 122,10–122,10 ($E_{\rm up}=92$ K), and the Band 3
image at 3″ resolution of CH3OCHO 100,10–90,9 ($E_{\rm up}=30$ K).

Around the frequency of H2CCO 31,13–121,12 ($E_{\rm up}=101$ K) line an
extended emission can be seen in the outflow, this is however coincident with
the NH2CHO which is only 4 km s-1 apart (Fig. LABEL:fig:fig9_extra). NH2CHO is
more commonly observed in the shocked regions than H2CCO (Ceccarelli et al.
2017; Codella et al. 2017); it is possible that those lines are blended. More
low-energy transitions would be required to confirm the identification of
these lines.

The velocities observed for ice-mantle tracers in the outflow are <15 km s-1
with respect to the systemic velocity. This is slower than the CO and SiO
outflow line wings which have velocities up to 20–30 km s-1. On the other
hand, the lines are clearly broader than those of molecular tracers of UV-
irradiated regions that trace passively heated gas (see Section 4.3).

Ice-mantle tracers are also detected in SMM3 and B1-c, two Class 0 protostars
(Fig. LABEL:fig:161 and LABEL:fig:151). For SMM3 only a lower $A_{ij}$
transition of CH3OH with $E_{\rm up}=$ 61 K was targeted and not detected
towards this source. B1-c has its methanol emission confused with the high-
velocity SO emission, whereas emission from other COMs in the outflow is weak
and appears only on the redshifted part of the outflow. Thus, S68N is the best
case to study the composition of shock-released ice mantles.

###  3.4 Outflow cavity walls

In this section, we highlight key molecules detected in the outflow cavity
walls. It is challenging to precisely distinguish cavity walls from the
outflowing material. The velocity of the gas in the cavity walls should be
lower than in the outflow, as the cavity wall contains envelope material at
rest but which is passively heated by UV radiation. We first discuss maps of
species associated with cavity walls and then their line profiles.

Figure 8: Spectra obtained at the cavity wall positions for hydrocarbons (red)
and ice-mantle tracers (blue). Left: C2H 32.5,2–21.5,1 ($E_{\rm up}=$ 25 K)
and CH3OH 21,0–10,1 ($E_{\rm up}=$ 28 K) spectra for S68N, Middle: C2H
32.5,2–21.5,1 ($E_{\rm up}=$ 25 K ) and SO 67–56 ($E_{\rm up}=48$ K) spectra
for B1-c, Right: c-C3H2 $4_{4,1}$–$3_{3,0}$ ($E_{\rm up}=32$ K), H2CO
32,1–22,0 ($E_{\rm up}=68$ K), and 13CS 5–4 ($E_{\rm up}=33$ K) spectra for
SMM3. Figure 9: Compact emission for B1-c and SMM3 for various molecules
tracing the warm inner envelope (hot core). Moment 0 maps shown in colorscale
integrated from -3 to 3 km s-1 w.r.t $\varv_{\rm sys}$. 1.3 mm continuum
presented in contours.

Figure 7 presents integrated emission maps of key tracers discussed in this
section for two examples: SMM3 and B1-c. Plots for the remaining sources
(S68N, Emb8N, and TMC1) are presented in the Appendix in Fig.
LABEL:fig:hydrocarbons_no2. Key tracers of the outflow cavity walls are simple
unsaturated hydrocarbon molecules: c-C3H2 $4_{4,1}$–$3_{3,0}$ ($E_{\rm up}=32$
K) for SMM3 and TMC1, and C2H $3_{2.5,3}$–$2_{1.5,2}$ ($E_{\rm up}=25$ K) for
S68N, B1-c, and Emb8N, as seen in maps obtained with ALMA Band 6 at
0$\aas@@fstack{\prime\prime}$5 resolution.

In Emb8N and SMM3 the emission from C2H and c-C3H2, respectively, is
symmetric; it appears similar in extent and shape on both sides of the
continuum source. Comparison with CO emission, which traces the bulk of the
outflowing gas, indicates that the hydrocarbons are located in the outflow
cavity walls close to the source. A higher energy transition of c-C3H2
72,6–71,7 $E_{\rm up}$=61 K is seen towards SMM3 closer to the protostar
compared with the lower-energy transition. This reflects the increase of
temperature of the cavity walls closer to the source. For B1-c, the emission
from C2H is U-shaped suggestive of a cavity wall, stronger on the blueshifted
side of the outflow, which could be either a projection effect or an asymmetry
in the envelope structure. For B1-c no ALMA CO observations exist to compare
with the bulk of the outflow at comparable resolution; however, the other
outflow tracer, SO, confirms the outflow direction and rough extent of the
outflow cavity walls. Moreover, the shape of the cavity walls is consistent
with the appearance of the CO 3–2 outflow observed with SMA toward this source
at 4″ resolution (Stephens et al. 2018).

S68N presents a chaotic structure (Fig. LABEL:fig:hydrocarbons_no2, top), but
C2H is found elongated in the outflow direction. While it is difficult to
identify the cavity wall, the C2H emission surrounds the CO outflow emission.
The C2H emission toward this source is asymmetric, with stronger emission in
the blueshifted part of the outflow. Emission of c-C3H2 toward the Class I
source TMC1 (Fig. LABEL:fig:hydrocarbons_no2, bottom) is not directly related
to the cavity walls, but is extended perpendicular to the outflow, which
suggests that c-C3H2 traces the envelope or extended disk material.

Figure 7 also shows CN 1–0 ($E_{\rm up}$ = 5 K) observed at 3″ resolution in
Band 3 for SMM3 and B1-c. Compared with C2H, CN is tracing similar regions. In
S68N CN has a similar extent as C2H but not over the full extent of the
outflow traced by CO. In B1-c, the CN emission has a similar shape of the
cavity wall cone as seen in C2H, but also a significant contribution from
larger scales is detected. In all cases the CN emission avoids the central
region, which likely results from on-source absorption by the foreground CN
molecules.

In some cases, like SMM3, the extent of the CN is broader than hydrocarbons,
more comparable with CO, but the narrow linewidths suggest that this emission
is still associated with passively irradiated envelope rather than with the
entrained outflow.

TMC1 presents a high-resolution example of CN emission (Fig.
LABEL:fig:hydrocarbons_no2). The offset between CO and CN reveals a physical
structure of the inner regions of the protostellar system: the entrained
outflow traced with CO appears closer to the jet axis, while CN highlights the
border between the outflow cavity wall and quiescent envelope. CN is sensitive
to UV radiation, as it can be produced with atomic C and N, whose abundances
are enhanced in PDRs, with UV photodissociation of HCN contributing as well
(Fuente et al. 1993; Jansen et al. 1995; Walsh et al. 2010; Visser et al.
2018).

H13CO+ emission is presented for B1c (Fig. 7), and S68N and Emb8N (Fig.
LABEL:fig:hydrocarbons_no2) observed in Band 6 at
0$\aas@@fstack{\prime\prime}$4 resolution. The bulk of the emission from this
molecule appears to be related to the cold envelope, however streams of
material can be seen in B1-c and S68N. The streams of gas observed in H13CO+
are coincident with the cavity wall observed in C2H. As H13CO+ is expected to
probe the dense envelope, the similarity of the morphology of the traced
material between H13CO+ and C2H and CN shows that the envelope material is UV-
irradiated. 13CS observed in Band 6 with the 12m array is detected for SMM3
(Fig. 7). The morphology of 13CS emission is very similar to that of c-C3H2.

In Fig. 8 spectra of ice mantle tracers, CH3OH and H2CO, and the hydrocarbon
molecules, C2H and C3H2 are presented. All spectra are shifted by their source
velocity to zero km s-1. In case of S68N, it is seen that C2H and CH3OH have
very similar line profiles indicating that they trace similar material. The
width of $\sim 10$ km-1 suggests that this material is entrained with the
outflow. A narrow component appears to be superposed at systemic velocities.
Note that fine splitting of C2H blends the spectra although the other
transition at +2 km s-1 does not affect the blueshifted velocity component. In
contrast, B1c shows only remarkably narrow C2H line profiles with a FWHM of
$\sim$ 2 km s-1, and SMM3 has similarly narrow c-C3H2 and 13CS lines compared
with broader H2CO emission (Fig. 8, right).

Figure 10: Images of Class I disks observed with ALMA 12m in Band 6 (van ’t
Hoff et al. 2020a). Top: Moment 0 maps of C17O at
0$\aas@@fstack{\prime\prime}$4 resolution integrated from -10 to 10 km s-1
w.r.t $\varv_{\rm sys}$. Bottom: Moment 0 maps of CN integrated from -2 to 2
km s-1 w.r.t $\varv_{\rm sys}$. Continuum contours in black.

###  3.5 Inner envelope

Fig. 9 shows emission from H2CCO, HNCO, t-HCOOH, SO, H2CS, and H13CN for the
Class 0 protostar B1-c, and H2S and OCS lines for SMM3; all observed with the
ALMA 12m array at 0$\aas@@fstack{\prime\prime}$5\. Several molecules tracing
the inner hot envelope are detected in 7m data and are presented in Fig.
LABEL:fig:102. CCS is detected for SMM3, BHR71, and IRAS 4B, while non-
detected for B1-c, TMC1, and Per-emb-25. OCS and H2S (Fig. LABEL:fig:102) for
all targeted in 7m Class 0 sources while non-detected for Class I sources.

The SO 67–56 $E_{\rm up}=48$ K line is seen to peak on the central source for
B1-c and S68N (Fig. LABEL:fig:so_outflow). SO has already been discussed in
the outflow (Section 4.2.2), but it is also prominent in the inner envelope.
While the spatial resolution does not allow to disentangle the hot core
emission from the small-scale outflow on a few hundred au scale, there is a
difference between these two sources and SMM3 and Emb8N. The latter two
sources show a substantial decrease in SO intensity towards the continuum
peak, i.e., they have prominent SO emission in the outflow, but not from the
hot core. This suggests that sources like B1-c and S68N, which are bright in
SO toward the continuum emission peak, have an additional component
responsible for SO emission. This is highlighted by the narrower lines of SO
toward the continuum peak compared with the outflow in S68N (Fig.
LABEL:fig:so_spectra_v1). The narrow component, visible in spectra taken on-
source, has a width of $\sim$ 5 km s-1. The main component of the spectrum
taken in the blueshifted outflow has a similar width but has a more prominent
line wing up to 20 km s-1.

TMC1 clearly shows SO emission toward both components of the binary system,
slightly offset from their peak positions (Fig. LABEL:fig:so_outflow). There
is also a molecular ridge present in SO close to the disk-envelope interface.

HNCO and HN13CO are detected toward B1-c and S68N peaking on source in higher
$E_{\rm up}$ transitions (Nazari et al. 2021). For lines with $E_{\rm up}<$90
K an extended component is also detected in the outflow. SMM3 and Emb8N have
no detections of HNCO on source, but for SMM3 this molecule appears in the
outflow. For all Class I sources where the relatively strong HNCO
110,11–100,10 line ($A_{ij}=2\times 10^{-4}$ s-1, $E_{\rm up}$=70 K) was
targeted, it was not detected.

OCS is detected toward SMM3 peaking in the center; S68N and SMM1 show
centrally peaked O13CS detection, a minor isotopologue signalling a high
abundance of OCS (Fig. LABEL:fig:o13cs_appendix). In all cases the emission is
moderately resolved; of size $\sim$ 200 au in case of SMM3 and detected up to
500 au away from source for S68N and SMM1.

H2S shows strong emission toward SMM3 and is also weakly present in TMC1. For
SMM3 the emission is resolved along the outflow direction and perpendicular to
the expected disk axis. Those are the only two sources for which H2S 12m array
data were taken. Additionally, the 7m data presented in the Appendix (Fig.
LABEL:fig:102) show prominent, centrally peaked H2S emission for four more
Class 0 sources.

H2CS is detected for B1-c, S68N and L1448-mm through a line with $E_{\rm up}=$
38 K. Another transition with $E_{\rm up}=46$ K is found in Class I disks:
IRAS-04302, L1489, and TMC1A. In B1-c, the emission is marginally resolved,
while in IRAS-04302 the molecule is clearly seen across the midplane,
indicating sublimation from icy grains at temperatures of at least 20 K (van
’t Hoff et al. 2020a; Podio et al. 2020).

H2CCO is detected for B1-c and S68N. For the lower $E_{\rm up}=100$ K
transition, the molecule is also detected in the outflow. For Class I sources,
the transition at comparable energy is not detected. HCOOH is detected for
B1-c and S68N with low-energy transitions that are seen both on source and in
the outflow, while the higher energy line ($E_{\rm up}=83$ K) is seen only on
source. H13CN is detected in B1-c and L1448-mm on source, additionally to the
outflow component.

Disks are commonly observed in Class I sources (Harsono et al. 2014; Yen et
al. 2017), as the envelope clears out. Fig. 10 presents maps of the C17O 2–1
($E_{\rm up}=$ 16 K) and CN 2–1 ($E_{\rm up}=$16 K) lines toward Class I disks
and L1527-IRS, which is identified as a Class 0/I object, observed with ALMA
12m at 0$\aas@@fstack{\prime\prime}$3 resolution. The C17O and H2CO emission
for disks in Taurus using these data are analyzed in detail by van ’t Hoff et
al. (2020a). Here we discuss CN in comparison with C17O. C17O is observed
concentrated towards the continuum emission for all disks, and is a much
cleaner tracer of the disk than any other more abundant CO isotopologues,
although even C17O still shows some trace emission from the surrounding
envelope.

Figure 11: Summary cartoon presenting key molecular tracers of different
components, limited to molecules presented in this work except for NH3 and
N2H+ which are important large scale envelope tracers. CN molecule is shown on
top of the disk as it traces the disk atmosphere and not the midplane.
Protostar indicated in the center in yellow.

##  4 Discussion

Figure 11 summarizes current understanding about which molecule traces which
component in Class 0/I protostars. Table 2 lists all molecules discussed in
this work with indication of the physical component that they trace. Below, we
will address several implications of the results. We note that the discussion
is confined to the molecules targeted and detected in our datasets, but is by
no means complete. Other relevant tracers not discussed here are NH3 and N2H+,
which are tracing the envelope well (e.g., Tobin et al. 2011; Chen et al.
2007). PO and PN have been recently suggested as a cavity walls or outflow
tracers, but have very low abundances (Bergner et al. 2019). SO2 is an
important tracer of the warm gas in the disk-envelope interface (Artur de la
Villarmois et al. 2019). Finally, this work focuses on low-energy transitions
observed with ALMA $E_{\rm up}<$100 K with the exception for hot core tracers.

Table 2: Molecular tracers of physical components exclusive to this work Molecule | Envelope | Disk | Hot | Outflow | Jet | Cavity
---|---|---|---|---|---|---
|  |  | core |  |  |
CO | ✓ | — | — | ✓ | ✓ | ✓
C18O | ✓ | ✓ | — | ✓ | — | ✓
13CO | ✓ | ✓ | — | ✓ | — | ✓
C17O | ✓ | ✓ | ✓ | — | — | ✓
H13CO+ | ✓ | — | — | ✓ | — | ✓
H2CO | ✓ | ✓ | ✓ | ✓ | ✓ | —
H2CS | — | ✓ | ✓ | — | — | —
SO | — | ✓ | ✓ | ✓ | ✓ | —
SiO | — | — | — | ✓ | ✓ | —
DCO+ | ✓ | — | — | — | — | —
N2D+ | ✓ | — | — | — | — | —
OCS | — | — | ✓ | — | — | —
O13CS | — | — | ✓ | — | — | —
HNCO | — | — | ✓ | ✓ | — | —
H2CCO | — | — | ✓ | ✓ | — | —
HCOOH | — | — | ✓ | ✓ | — | —
13CS | ✓ | — | — | — | — | ✓
CN | ✓ | ✓ | — | — | — | ✓
HCN | — | — | — | ✓ | — | —
H13CN | — | — | ✓ | ✓ | — | —
C2H | ✓ | — | — | ✓ | — | ✓
c-C3H2 | — | — | — | — | — | ✓
CH3CHO | — | — | ✓ | ✓ | — | —
CH3CN | — | — | ✓ | ✓ | — | —
CH3OCHO | — | — | ✓ | ✓ | — | —
CH3OH | — | — | ✓ | ✓ | — | —
13CH3OH | — | — | ✓ | — | — | —

###  4.1 Cold envelope

####  4.1.1 C18O

C18O is a good tracer of high column density material because of its low
critical density. However, it becomes less abundant as soon as the dust
temperature drops below the CO freeze-out temperature ($\sim$ 20–25 K). There
is also a density threshold: CO freeze-out only occurs at densities above
$\sim 10^{4}$–$10^{5}$ cm-3, because at lower densities the timescales for
freeze-out are longer than the lifetime of the core (Caselli et al. 1999;
Jørgensen et al. 2005b).

For 4 out of 6 sources presented here at 6″ resolution, Kristensen et al.
(2012) performed modeling of the SED and sub-mm spatial extent using the DUSTY
code (Ivezic & Elitzur 1997). The results provide, among other properties, a
temperature structure throughout the envelope and the radius at which the
temperature drops below 10 K, which is considered as the border between
envelope and the parent cloud. Kristensen et al. (2012) obtained radii of
3800, 5000, 6700, and 9900 au for IRAS4B, TMC1, SMM3, and BHR71, respectively.
The compact C18O emission observed on-source and its non-detection over the
full expected extent of the protostellar envelope can be explained by CO
freeze-out occurring already within the inner 1800 – 2500 au radius, which is
the spatial resolution of our observations for Class 0 sources. This upper
limit on the CO snowline is consistent with CO snowlines typically observed
and modeled toward other Class 0 protostars (Jørgensen et al. 2004; Anderl et
al. 2016; Hsieh et al. 2019).

Equation 1 from Frimann et al. (2017) allows to calculate the expected CO snow
line for the current luminosity of the example sources B1-c and TMC1 presented
in Fig. 2 in the absence of an outburst. For B1-c the snowline is expected to
be at 200–400 au radius depending on the assumptions of the sublimation
temperature (larger radii for 21 K and smaller for 28 K), while the TMC1 CO
snowline is expected to be at 100–200 au. For the most luminous source with
C18O 7m observations available, the expected radius is at 400–750 au.
Therefore clearly in all cases the expected CO snowline is well within the 7m
beam. In high-resolution studies, the CO emission is often seen at greater
distances than expected from the current luminosities of those protostars.
This is attributed to accretion bursts of material which increase their
luminosities resulting in a shift of the observed CO emission radius up to a
few times its expected value (but usually still within a 1000 au radius)
(Jørgensen et al. 2015; Frimann et al. 2017; Hsieh et al. 2019).

####  4.1.2 N2D+ and DCO+

Both DCO+ and N2D+ are considered cold gas tracers (e.g., Qi et al. 2015).
N2D+ is efficiently destroyed by CO in the gas-phase, therefore freeze-out of
CO results in N2D+ being retained in the gas-phase at larger radii of the
envelope, where temperatures are lower. This behaviour has been demonstrated
in several other protostellar sources by Tobin et al. (2011, 2013, 2019).

Both DCO+ and N2D+ are produced through reactions with H2D+. At cold
temperatures the H2D+ abundance is enhanced through the H${}_{3}^{+}$ \+ HD
$\rightarrow$ H2D+ \+ H2 reaction, which is exothermic by 230 K. As the
reverse reaction is endothermic, low temperatures increase H2D+. Additionally,
both H${}_{3}^{+}$ and H2D+ are enhanced in gas where CO has been depleted.
However, the CO molecule is still needed for the production of DCO+ through
the H2D+ \+ CO reaction. Therefore, DCO+ is expected to be most abundant
around the CO snowline (Jørgensen et al. 2004; Mathews et al. 2013). Warmer
production routes through CH2D+ \+ CO are also possible (Wootten 1987; Favre
et al. 2015; Carney et al. 2018).

The difference between DCO+ and N2D+ chemistry is reflected in morphology of
both molecules in the dense regions close to the continuum peak. As DCO+
requires gas-phase CO for its formation, it peaks close to the CO snowline,
which is within the resolution of our observations ($\sim$1800–2500 au
radius), while N2D+ is only located where CO is not present in the gas phase.
Therefore we observe a significant decrease of N2D+ in the inner envelope. If
the warm production of DCO+ is triggered in the inner regions, this will
additionally produce DCO+ within the beam of our observations, hence DCO+ does
not decrease in the inner envelope. The extent of the DCO+ and N2D+ emission
in each source is comparable, ranging from $\sim$ 5000 au in B1-c and BHR71 to
1500 au in TMC1, suggesting that their outside radii trace the region where CO
becomes present again in the gas phase due to the low density.

The morphology of the emission from cold gas tracers such as DCO+ and N2D+ is
sensitive to the density and temperature profile of the system, which can be
affected by system geometry (i.e., outflow opening angle, disk flaring angle,
flattening of the envelope). DCO+ has been shown to increase its abundance in
the cold shadows of a large embedded disk (Murillo et al. 2015). Emission from
DCO+ and N2D$+$ is consistent with a picture of a dissipating envelope in
Class I sources, resulting in less dense, warmer gas surrounding the
protostar. TMC1 has an order of magnitude lower envelope mass compared to the
Class 0 sources (Table 1; Kristensen et al. 2012). This causes the extent of
the cold and dense region to shrink, preventing N2D+ from being detected, and
limiting the extent of DCO+ emission. The dense gas toward TMC1 is clearly
present only in the flattened structure surrounding the binary system, likely
forming a young, embedded disk. In fact this source is suggested to have a
rotationally-supported circumbinary disk (Harsono et al. 2014; van ’t Hoff et
al. 2020a). The geometry of the disk can create favourable conditions for the
DCO+ enhancement in the cold shadows of the disk.

Other relevant molecules that trace the quiescent envelope material but are
not presented here are HCO+ and H13CO+ (Hogerheijde et al. 1997; Jørgensen et
al. 2007; Hsieh et al. 2019, van ’t Hoff et al. in prep.). These molecules
have been shown to probe the material outside of the water snowline (Jørgensen
et al. 2013; van ’t Hoff et al. 2018a). As water sublimates at temperatures
$\sim 100$ K, much higher than CO, HCO+ can be seen throughout the envelope,
except for the warmest inner regions. N2H+ is tracing the envelope material
and CO snowline and has been shown to peak closer to the central protostar
than N2D+ (Tobin et al. 2013). Their ratio can potentially be used as an
evolutionary tracer of protostars (Emprechtinger et al. 2009). NH3 is also
observed to map similar components as N2H+ (Tobin et al. 2019).

In summary, the quiescent envelope material is traced by dense and/or cold gas
tracers. Chemical interactions result in N2D+ tracing the outer envelope where
CO is frozen-out, whereas DCO+ is seen both in the outer envelope as well as
in the inner regions, tracing the unresolved CO snowline. C18O is a good
tracer of dense ($n$ $>10^{5}$ cm-3) and warm ( $T$ $>$ 30 K) regions in the
inner 2000 au radius of the protostellar systems. The evolution of the
physical conditions from Class 0 to Class I is evident as the envelope becomes
less dense and the protostellar luminosity can heat up dust and gas more
easily.

###  4.2 Outflows and jets

Outflowing material from protostellar systems is best analyzed with kinematic
information. In the following section, we will discuss three different
componfents of protostellar outflows observed with the ALMA 12m array: 1) the
high-velocity jet (>30 km s-1), 2) the low-velocity outflow (<30 km s-1), 3)
the gas that results from the interaction with the outflow – ice sputtering
products at velocities close to that of ambient material, but with linewidths
significantly broader (up to 15 km s-1) than the quiescent envelope tracers.
The three components are presented in Fig. 12 with examples of S68N, a
representative case of a prominent outflow and ice sputtering, and L1448-mm, a
source with a prototypical high-velocity jet.

####  4.2.1 Extremely high-velocity (EHV) jet

There are 6 out of 7 Class 0 sources targeted at high-resolution strengthen
the conclusion that EHV jets are more common than previously thought in Class
0 sources (Podio et al. 2021). It is argued that the molecular jet tracers
have a very different physical origin than the protostellar outflows; contrary
to the low-velocity outflow, which consists mostly of entrained envelope
material, the EHV jet is expected to be directly launched from the innermost
region of the system (Tafalla et al. 2010; Lee 2020). The high-velocity jet is
comprised of atomic material which readily forms molecules in the high-density
clumps (internal working surfaces; Raga et al. 1990; Santiago-García et al.
2009) that are resulting from shocks in the jet, which are produced by the
velocity variations of the ejection. This in turn means that by observing the
high-velocity bullets, one gains insight on the variability of the accretion
process (Raga et al. 1990; Stone & Norman 1993). The new EHV sources B1c and
SMM3 show bullet spacings of 1200 au in B1c and 3200 au in SMM3, which can be
converted using the terminal velocity of the jets (not corrected for
inclination) to the dynamical ages of 80 and 250 years, respectively. If the
central mass of the protostar can be estimated this can be used to provide the
orbital period of the component causing the variability (Lee 2020).

The fact that the EHV jet tracers are dominated by O-bearing molecules has
been associated with a low C/O ratio in the jet material (Tafalla et al.
2010). For high mass-loss rates, molecules are produced efficiently in the jet
from the launched atomic material (Glassgold et al. 1991; Raga et al. 2005;
Tabone et al. 2020). Additionally the ratio of SiO-to-CO can indicate the
presence of dust in the launched material, which can in turn inform about the
jet launching radius, i.e., whether it is inside or outside the dust
sublimation radius (Tabone et al. 2020). The new detections of high-velocity
jets suggest that this process may be occurring in every young Class 0 object.
Studying large samples of objects with ALMA and combining with multi-
transition observations can unveil the atomic abundances of the inner regions,
which are difficult to measure otherwise directly (McClure 2019).

The presence of H2CO in the jet could result from gas-phase production through
the reaction of CH3 with O. Alternatively, if the icy grains were launched
with the jet, they could be sputtered in the internal working surfaces at high
velocities (Tychoniec et al. 2019).

In summary: O-bearing species such as CO, SiO, SO, and H2CO observed at high
velocities are excellent tracers of the chemistry within the protostellar jet.
Those molecules most likely formed in the internal working surfaces from the
material carried away from the launching region of the jet.

Figure 12: Maps of three different components of the outflow. Moment 0 maps
are presented in color scale with continuum emission at 1.3 mm in black
contours, both obtained with ALMA 12m observations. Left: An extremely high-
velocity (EHV) molecular jet illustrated with the SiO (4–3) map for L1448-mm
integrated from -70 to -50 and from 50 to 70 km s-1 w.r.t $\varv_{\rm sys}$.
Middle: Low-velocity outflow illustrated with the CO 2–1 map for S68N
integrated from -15 to -3 and from 3 to 15 km s-1 w.r.t $\varv_{\rm sys}$.
Right: Ice mantle content released with shock sputtering presented with the
CH3OH (21,0 – 10,0) map for S68N integrated from -8 to -1 and from 1 to 8 km
s-1 w.r.t $\varv_{\rm sys}$. The CO outflow from Ser-emb-8N is present at the
edge of the map.

####  4.2.2 Low-velocity outflow

The CO molecule traces the bulk of the gas as it is the most abundant molecule
detectable in the sub-mm regime. It serves as an indicator of the outflow
extent and gas morphology, as it is not affected by chemical processing in
shocks. It can also be used to quantify the total mass-loss rates, using the
dense cloud abundance ratio of CO/H2 $\sim$ 10-4. A well-known correlation
between molecular outflows and protostellar luminosities indicates a strong
link between the accretion and ejection processes (Cabrit & Bertout 1992;
Bontemps et al. 1996; Mottram et al. 2017). The decrease in accretion and
total envelope mass with evolution of the system also results in fainter, less
powerful outflows. The weak emission from the Class I source TMC1, which is an
order of magnitude lower in intensity with respect to outflows from Class 0
protostars, is consistent with this trend.

SiO is a molecule that is enhanced by several orders of magnitude in shocks
compared with gas in cold and dense clouds where most of the Si is locked in
the grains (e.g., Guilloteau et al. 1992; Dutrey et al. 1997). Shocks release
Si atoms from the grains by means of sputtering and grain destruction, leading
to subsequent reactions with OH, another product of shocks, forming SiO
(Caselli et al. 1997; Schilke et al. 1997; Gusdorf et al. 2008b, a). Thus, SiO
is much more prominent in high-velocity gas, where grains are more efficiently
destroyed (see Section 4.2.1), than in low-velocity shocks.

SO is enhanced in shocks through reactions of atomic S released from the
grains with OH, as well as through H2S converted to SO with atomic oxygen and
OH (Hartquist et al. 1980; Millar & Williams 1993). Shocks could also explain
the emission toward TMC1, where weak accretion shocks onto the disk could
enhance the SO abundances (Sakai et al. 2014b; Yen et al. 2014; Podio et al.
2015). Overall, there is a clear decrease of both SiO and SO low-velocity
emission from Class 0 to Class I. Either the less powerful jet cannot destroy
the grains and create conditions for the production of SiO and SO and/or the
much less dense envelope and outflow cavity walls do not provide enough dust
grains for creating large column densities of those molecules; additionally,
the excitation conditions might change significantly with evolution of the
protostellar system, hampering the detection of even low-$J$ SO and SiO
transitions with critical densities in the $10^{5}-10^{6}$ cm-3 range.

HCN is associated with the most energetic outflows (Jørgensen et al. 2004;
Walker-Smith et al. 2014) and is enhanced at high temperatures in shocks.
H13CN is likely associated with both the hot inner regions (not detected in
HCN due to optical thickness of the line) and the low-velocity outflow
(Tychoniec et al. 2019; Yang et al. 2020). The geometry of H13CN emission seen
in HH211 resembles a cavity wall: it could be a result of CN produced by UV-
photodissociation and subsequent production of HCN via the H2 \+ CN reaction,
which requires high temperatures (Bruderer et al. 2009; Visser et al. 2018).
The fact that H13CN is seen at outflow velocities shows that shocks are
required to produce HCN, which is likely released at the cavity walls and then
dragged with the outflow.

####  4.2.3 Shock sputtering products

Ice-mantle tracers are detected in outflows of SMM3, B1-c and S68N protostars.
As indicated in Section 3.3 S68N is the best test-case to study the
composition of shock-released ice mantles. The morphology of the emission of
ice-mantle tracers in S68N is somewhat uniform. CH3CN and HNCO are clearly
brighter on the redshifted part of the outflow (south-east) and CH3CHO are
brighter in the blueshifted (north-west) side. CH3OH and H2CO show an even
distribution between the two lobes. The peak intensity for all species occurs
at significant distances from the source ($\sim$ 5000 au) and in some cases
the emission drops below the detection limit closer to the source. This is
contrary to the CO emission, which can be traced all the way back to the
central source. In all tracers the emission is also detected at the continuum
position, however, this emission has a narrow profile and results from thermal
sublimation of ices in the hot core of S68N (van Gelder et al. 2020).

The velocities observed for CH3OH and other ice-mantle tracers are high enough
for these molecules to be material near the outflow cavity walls, where ice
mantles could be sputtered. These molecules therefore most likely trace low-
velocity entrained material with a considerable population of ice-coated
grains that are sputtered in the shock (Tielens et al. 1994; Buckle & Fuller
2002; Arce et al. 2008; Jiménez-Serra et al. 2008; Burkhardt et al. 2016).
Thermal desorption of molecules from grains is not likely as dust temperatures
at distances of few times 103 au from the protostar are well below the
desorption temperatures for most COMs. The fact that there is no enhancement
of these tracers closer to the source also argues against emission being
related to high temperature.

###  4.3 Outflow cavity walls

Emission from hydrocarbons and CN in the outflow cavity walls is directly
related to the exposure of those regions to the UV radiation from the
accreting protostar. Both c-C3H2 and C2H have been prominently observed in
PDRs such as the Horsehead Nebula and the Orion Bar tracing the layers of the
cloud where UV-radiation photodissociates molecules, which helps to maintain
high atomic carbon abundance in the gas-phase that is needed to build these
molecules. C2H is enhanced in the presence of UV radiation at cloud densities
(Fuente et al. 1993; Hogerheijde et al. 1995) and c-C3H2 usually shows a good
correlation with C2H (Teyssier et al. 2004). Both molecules have efficient
formation routes involving C and C+, although models with only PDR chemistry
tend to underpredict their abundances, especially for c-C3H2. A proposed
additional mechanism is the top-down destruction of PAHs (Teyssier et al.
2004; Pety et al. 2005, 2012; Guzmán et al. 2015); the spatial coincidence of
PAH emission bands with hydrocarbons in PDRs is consistent with that
interpretation (van der Wiel et al. 2009).

It is instructive to compare the conditions between classic PDRs and outflow
cavity walls around low-mass protostars.The $G_{0}$ value for Orion Bar is
estimated at 2.6 $\times$ 104 (Marconi et al. 1998), while the Horsehead
Nebula has a much more moderate radiation field of $\leavevmode\nobreak\
10^{2}$ (Abergel et al. 2003). The UV radiation field around low-mass
protostars measured by various tracers is 102–103 at $\sim$ 1000 au from the
protostar (Benz et al. 2016; Yıldız et al. 2015; Karska et al. 2018),
therefore the PDR origin of small hydrocarbons is plausible. The top-down
production of hydrocarbons due to PAHs destruction does not appear to be an
efficient route here, as PAHs are not commonly observed in low-mass
protostellar systems (Geers et al. 2009), and the UV-fields required for this
process are above 103 (Abergel et al. 2003).

The difference in morphology of hydrocarbons between Class 0 and Class I
systems – outflow cavity walls in Class 0 versus rotating disk-like structure
in Class I – is most likely related to the evolution of the protostellar
systems. Class 0 sources have a dense envelope and the UV radiation can only
penetrate the exposed outflow cavity walls, while for Class I it is likely
much easier for both the UV radiation from the accreting protostar and the
interstellar radiation to reach deeper into the envelope or disk. This is
consistent with emission from well-defined cavity walls seen toward those
sources.

The case of moderate $\sim 10$ km s-1 velocity material observed toward S68N
indicates that the C2H line does not in all sources trace exclusively the
quiescent cavity walls. The profile is consistent with the observed morphology
of the line (see Fig. LABEL:fig:hydrocarbons_no2) – emission is seen up to a
few thousand au from the source and its shape does not resemble a cavity wall
as clearly as in other sources. The narrow component centered at systemic
velocity seen in Fig. 8 indicates that while the broad component might be
dominating the emission, the UV-irradiated cavity wall also contributes to the
emission observed for S68N.

S68N could be a very young source, as the chaotic structure of its outflow and
envelope indicates (Le Gouellec et al. 2019). The high abundances of freshly
released ice-mantle components described in Section 4.2.3 are consistent with
this interpretation. It is also possible that UV radiation produced locally in
shocks is causing the enhancement of C2H emission at higher velocities.

The 13CS molecule, as a high-density tracer likely traces the material piling
up on the cavity walls pushed by the outflow. This emission is usually slow
($\pm$ 2 km s-1 ) indicating that this is not outflowing gas but rather
envelope material on the outflowing cavity walls. The non-detection of 13CS in
12m data (Fig. LABEL:fig:hydrocarbons_no2) towards the Class I source TMC1 is
consistent with the dissipating envelope as the source evolves, hence no high-
density material is seen in the remnant cavity walls, even though they are
still highlighted by the CO emission. In the remaining Class 0 sources with
the 13CS detections – Emb8N, SMM1 and S68N (Fig. LABEL:fig:fig_E2), – emission
also follows other cavity wall tracers. SMM1 is the only source that shows
higher velocity structure (¿ 7 km s-1) of the 13CS emission on one side of the
cavity wall. As this is spatially coincident with the high-velocity jet
observed in CO (Hull et al. 2016), this emission might be related to the
material released with the jet from the envelope.

To summarize, we observe the hydrocarbons C2H and c-C3H2, as well as CN,
H13CO+ and 13CS in the outflow cavity walls of Class 0 protostars. C2H, as the
most abundant of the hydrocarbons presented here, is also seen prominently
across the envelope at velocities comparable to the low-velocity outflow
whereas c-C3H2 appears as a clean tracer of the quiescent, UV-irradiated gas
in the cavity walls in the Class 0 sources.

###  4.4 Warm inner envelope

####  4.4.1 Compact emission

The inner regions of young protostellar systems are characterized by high
temperatures which result in a rich chemistry as molecules that form
efficiently in ices on grains in cold clouds sublimate into the gas phase.
Complex organic molecules (COMs) detected in these datasets are discussed
quantitatively in detail elsewhere, for both O-bearing by van Gelder et al.
(2020) and for N-bearing species by (Nazari et al. 2021) In this section we
focus on smaller molecules that also trace the innermost hot core regions and
therefore are likely abundant in ices. This includes several small S-bearing
molecules.

Both B1-c and S68N are sources characterized as hot cores (Bergner et al.
2017; van Gelder et al. 2020), which means that the conditions in their inner
regions are favourable for release of molecules from the ice mantles. For the
well-studied case of IRAS16293-2422, Drozdovskaya et al. (2018) have also
identified a hot core component of SO based on isotopologue data, in addition
to an SO component in the large scale outflow. Overall, SO observed in
protostars appears to be related mostly to evaporation of grain mantle.

The SO appearance close to the central protostar could be related to accretion
shocks onto the disk, which are weaker than shocks that cause the SO emission
seen in the outflows (Section 4); in accretion shocks, SO can be released from
the icy mantles with the infalling material (Sakai et al. 2014b). However,
narrow linewidths of SO toward TMC1 seem to rule out the accretion shock
scenario, and points to emission along the cavity walls (Harsono et al. 2021).

HNCO emission has been modeled by Hernández-Gómez et al. (2018) and suggested
to be a superposition of both warm inner regions of the envelope as well as
the colder, outer envelope. Its similar behaviour to sulphur-bearing species,
also observed in our work, is proposed to be related to the fact that O2 and
OH are involved in formation of species like SO and HNCO (Rodríguez-Fernández
et al. 2010).

H2S is expected to be the dominant sulphur carrier in ices (Taquet et al.
2020). However, it has not yet been detected in ice absorption spectra to date
(Boogert et al. 2015). The weak emission from H2S in dark clouds has been
modeled as a result of the photodesorption of ices at the outside of the cloud
(like in the case of H2O, Caselli et al. 2012), while chemical desorption is
important for grains deeper inside the cloud but outside the water snowline
(Navarro-Almaida et al. 2020). These models are consistent with H2S ice
containing most of the sulphur. Multiple lines of H2CS are a powerful tool to
probe the warm ¿ 100 K, innermost regions of the protostellar systems (van ’t
Hoff et al. 2020b).

While B1-c and S68N are characterized as hot cores with many COM lines
detected (van Gelder et al. 2020), and L1448-mm has warm water in the inner
regions (Codella et al. 2010), SMM3 does not appear to have significant
emission from COMs. Moreover, simple molecules associated with the hot cores
for B1-c, such as SO and H2CO, are only seen outside of the SMM3 central
source. While it is possible that the optically thick continuum prevents a
detection of COMs in its inner envelope (De Simone et al. 2020a), differences
in chemistry or physical structure (e.g., a large cold disk, see Section 3.5)
between the SMM3 and the hot core sources are also possible. The fact that
emission from OCS and H2S is centrally peaked (Fig. 9, bottom row) suggests
that continuum optical depth is not an issue, although both those species
could be a result of grain destruction or ices sputtering, therefore not
necessarily coming from the midplane but rather from the surface of a disk-
like structure. The additional detection of CCS in 7 data towards SMM3, which
is on the other hand not detected in B1-c could hint at different chemical
composition of the two protostellar systems.

####  4.4.2 Embedded disks

In the case of large edge-on disks like IRAS-04302 and L1527, the vertical
structure of the emission can be probed, as well as the radius where CO
freeze-out occurs. Overall, Class I disks are warmer than their Class II
counterparts with CO freeze-out taking place only in the outermost regions
(Harsono et al. 2015a; van ’t Hoff et al. 2018a, 2020a; Zhang et al. 2020).

We detect CN in all observed Class I disks, although it does not trace the
midplane of the disk. In the near edge-on example of IRAS-04302, the CN
emission originates from the upper layers of the disks, in the same direction
as the outflow, which is perpendicular to the disk in this source. This opens
a possibility that the emission is also related to the irradiated residual
cavity walls in those sources. TMC1A is a clear example where CN is tracing
the same material as probed by Bjerkeli et al. (2016) in CO, which is
attributed to a disk wind. TMC1 and L1527 show CN oriented in the same
direction as the disk; in TMC1 there is also a clear filament structure on
larger scales irradiated by the UV from the protostar, seen also in other
tracers.

In comparison with the C17O emission, which traces the midplane disk, CN thus
appears in the upper layers and in the outflow, therefore in most cases the
two molecules are mutually exclusive. This picture is consistent with the bulk
density traced by CO and the irradiated layers of the disk and envelope
exposed to UV traced by CN. Recent observations of a sample of Class I sources
including IRAS-04302 by Garufi et al. (2021) is consistent with CN not tracing
the disk midplane.

In the younger Class 0 or borderline Class 0/I sources characterization of the
disk is much more difficult because of the strong envelope emission.
Nevertheless, several Keplerian disks have been identified with observations
of CO isotopologues like C18O and 13CO (Tobin et al. 2012; Murillo et al.
2013). Our data allow us to investigate this for the case of SMM3. Fig.
LABEL:fig:smm3_disk shows the red and blue-shifted emission from C18O toward
SMM3. There is a clear rotational signature in the direction perpendicular to
the outflow on scales of a few hundred au. However, to unambiguously identify
the disk and its radius, higher spatial and spectral resolution data are
necessary.

Figure 13: Ratio of complex organic molecules with respect to CH3OH in S68N.
The ratios are shown for a 3″ region in both redshifted and blueshifted part
of the outflow. The abundance ratios in the hot cores are taken from van
Gelder et al. (2020) and Nazari et al. (2021). The errorbars mark the lower
and upper limit on the ratios, as calculated from the optically thick CH3OH
emission (lower limit) and 13CH3OH (upper limit). All outflow abundances
obtained from the Band 6 data at 0$\aas@@fstack{\prime\prime}$45 resolution
except for CH3OCHO and CH3CN which were derived from Band 3 data at 3″.

###  4.5 COMs in outflow versus hot core

The complex organic molecules are already detected in the prestellar stage of
star formation (Bacmann et al. 2012; Scibelli & Shirley 2020), where they are
efficiently produced on the surfaces of icy grains (e.g., Watanabe et al.
2004; Öberg et al. 2009). COMs can be subsequently released back into the gas
through various non-thermal desorption processes and/or be re-formed by gas-
phase reactions. In the inner envelope, temperatures are high enough that
thermal desorption of ices is enabled. In the outflow cavity walls, sputtering
by shocks can release ices from the grains. In the hot core, COMs are detected
through high excitation lines, while in the outflow, they are observed
primarily through transitions with low $E_{\rm up}$. Therefore, comparing the
observations of molecular complexity in outflows with hot cores can unveil if
there is any warm temperature processing of the ices and whether some
molecules also have a gas-phase production route.

Therefore, we compare abundance ratios between different ice mantle species
and methanol at three positions for S68N: one on the source, obtained from van
Gelder et al. (2020) and Nazari et al. (2021), and one each in the blueshifted
and redshifted part of the outflow, which are calculated from the fluxes
obtained in this work. We measure the abundance of the species in a 3″ region
centered on CH3OH peak on the blueshifted and redshifted regions. The size of
the region is based on the spatial resolution of the Band 3 data. The regions
are indicated in Fig. 6. From these regions, we extracted spectra and
calculated column densities using the spectral analysis tool CASSIS
111http://cassis.irap.omp.eu/.

For the CASSIS model, we assume $T_{\rm ex}$ = 20 K, typical of subthermally
excited molecules with large dipole moments in outflow gas. The FWHM of the
lines was fixed at 3 km s-1, which is the width of the CH3OH line. With those
parameters, an LTE calculation provides column densities for the observed line
intensities under the assumption of optically thin emission (see Tychoniec et
al. (2019) for a discussion on uncertainties in this method.)

Only a single CH3OH line 21,0–10,1 ($E_{\rm up}$ = 28 K) is detected in the
outflow of S68N. Escape probability calculations with RADEX (van der Tak et
al. 2007) show that even at the lower densities in outflows this CH3OH line is
likely optically thick. No optically thin CH3OH isotopologue is detected in
the outflow; therefore, we provide an upper limit for the 13CH3OH 21,1-10,1
($E_{\rm up}=$28 K). The upper limit on this 13CH3OH line of $1\times 10^{14}$
cm-2 translates to an upper limit on the CH3OH column density of $7\times
10^{15}$ cm-2 assuming 12C / 13C = 70 (Milam et al. 2005).

The column density calculated from the flux measurement of the CH3OH line
gives of $2\times 10^{15}$ cm-2; thus, the methanol abundance can be
underestimated by a factor of 4. With the information available, we assume
that the CH3OH column density is between $2\times 10^{15}$ cm-2 and $7\times
10^{15}$ cm-2 and we compare abundances of other molecules with respect to
CH3OH using this range of values. The hot core methanol column density of S68N
in van Gelder et al. (2020) is corrected for optical thickness with
CH${}_{3}^{18}$OH.

Figure 13 compares the abundance ratios of various species with respect to
methanol for S68N. First, it shows that the relative abundances are remarkably
similar on both sides of the outflow, well within the uncertainties. Second,
most abundances relative to CH3OH are found to be comparable to the S68N hot
core within our uncertainties and those of van Gelder et al. (2020) and Nazari
et al. (2021), noting that due to lack of optically thin CH3OH line detected
in the outflow, our uncertainties are larger. The greatest difference is seen
for CH3CHO and H2CCO, which could be attributed to the additional gas-phase
formation in shocks. While the difference is well beyond the conservative
uncertainties assumed here, this result should be confirmed with a larger
sample of sources with different properties. The jump in CH3CHO abundance with
respect to methanol in the outflow has been also reported for two other
sources - L1157 and IRAS4A (De Simone et al. 2020b).

For the well-studied L1157 outflow, Codella et al. (2020) find CH3CHO/CH3OH is
0.5%. which they find in agreement with ratios in hot cores toward different
protostars, but not including L1157 itself. In the case of S68N we uniquely
show a comparison between the emission from molecules in the outflows and hot
core for the same source. While modeling by Codella et al. (2020) shows that
the similar spatial origin of CH3OH and CH3CHO does not imply that they are
both solely grain sputtering products, an agreement between hot cores and
outflows could mean that similar processes are responsible for emission in
both regions.

It is interesting that the SMM3 outflow shows ice mantle tracers in the
outflow but not on the source. Thus, the lack of hot core emission is likely
due to the physical conditions in the inner regions such as a large disk
(Section 3.5) or continuum optical depth (De Simone et al. 2020a) and not to
the lack of complex molecules on the grains. Since methanol is not detected,
we cannot provide abundance ratios for this outflow. B1-c - another source
with an outflow containing ice mantle products, has methanol lines overlapping
with the high-velocity SO outflow, therefore precise abundance measurements
are not possible.

JWST will be able to provide information on the ice content due to rich
absorption spectra in the mid-IR. The abundances of COMs in the gas-phase
provided by ALMA (van Gelder et al. 2020; Nazari et al. 2021) can then be
directly compared with the ice content. So far, the ice content observed on
cloud scales with Spitzer does not show a correlation with the gas content but
those data do not probe thermally desorbed ices (Perotti et al. 2020).

###  4.6 Carbon-chains and other hydrocarbons vs COMs

Early observations at low spatial resolution suggested that large carbon-
chains such as C5H, HC7N, and HC9N and complex organic molecules are mutually
exclusive, and therefore the two were proposed to be tracers of two different
categories of protostellar sources driven by different chemistry, namely warm
carbon-chain chemistry (WCCC) and hot core chemistry (Sakai & Yamamoto 2013).
Warm carbon-chain chemistry is thought to occur above the sublimation
temperature of CH4 $\sim$ at 30 K. The proposed difference then lies in the
WCCC sources collapsing more rapidly than the hot core sources, which prevents
CO to accumulate on the ices, leaving a higher CH4 ice content. Thus, in WCCC
sources, there would be less CO to form more complex organic molecules such as
CH3OH, and an underabundance of COMs is indeed observed. This scenario has
been supported by the lower deuteration observed in WCCC sources (Sakai et al.
2009). With ALMA observations at higher resolution and sensitivity, sources
harboring both COMs and small hydrocarbons, c-C3H2, and C2H, have now been
observed (Imai et al. 2016; Oya et al. 2017).

Figure 14: Emission maps of an example COM and hydrocarbon toward B1-c. Left:
CH3CHO moment 0 map, Right: C2H moment 0 map; both integrated from -2 to 2 km
s-1 w.r.t $\varv_{\rm sys}$.

Our data also show detection of small hydrocarbons and COMs in the same
sources (Fig. 14). However, our high resolution maps now show those
hydrocarbons to originate in the UV irradiated outflow cavity walls (Section
6). Both c-C3H2 and C2H can be solely explained by PDR chemistry (Guzmán et
al. 2015; Cuadrado et al. 2015) even in moderate UV fields of low-mass
protostars and therefore WCCC chemistry does not need to be invoked for the
sources presented in this work. The PDR origin is consistent with our observed
geometry at the edges of the cavities, where lower extinction allows the UV
radiation from the protostellar accretion to easily reach this region. It also
explains the confinement of this emission to the inner 2000 au of the envelope
as the UV radiation decreases with distance from the source. A deep search for
larger hydrocarbons like HC5N in cavity walls could help to understand their
origin.

###  4.7 Class 0 vs Class I in different components

Many processes take place within the protostellar lifetime of a few $10^{5}$
yr that alter the chemical composition and physical conditions. First, the
envelope is dissipating, and the radius where key molecules are in the gas-
phase is becoming smaller (Jørgensen et al. 2005b). With high-resolution
studies, it is possible to probe different snowlines in unprecedented detail
now (van ’t Hoff et al. 2018a; Hsieh et al. 2019). There is a clear
evolutionary trend observed in temperature – Class 0 disks are generally
warmer than Class II disks, with Class I sources in between, which has major
impact on the richness of their spectra (van ’t Hoff et al. 2020a, b). This is
also observed in our sample where all Class 0 show either complex organic
molecules or small molecules related to grain evaporation but no evidence of
those molecules is found in Class I sources (Section 4.4.1). Table 3
summarizes the evolution of the molecular composition of protostellar systems.

Complex organic molecules are a key tracer of chemistry at the early stages.
We see them in the outflows only in Class 0 sources, likely due to the
decrease of the outflow force and mass and envelope density. In the inner
regions, physical conditions such as temperature and density also favour the
early stages for the abundances of COMs. In rare cases of rich chemistry in
Class I disks, it is usually attributed to accretion outbursts heating up the
disk (van ’t Hoff et al. 2018c; Lee et al. 2019a).

Molecular EHV jets are only present in the Class 0 phase (Nisini et al. 2015).
In our sample 6 out of 7 sources have high-velocity component and none is
detected in Class I (Section 4.2.1), whereas the lower velocity outflows in
these stages are rich in ice mantle tracers (Section 4.2.3). The Class I jets
are invisible in molecular tracers, and their outflows show CO primarily, with
much less contribution from shock tracers such as SO and SiO (Section 4.2.2).

Cavity walls are seen prominently in the Class 0 sources in molecules whose
abundances are sensitive to UV radiation. In Class I sources, a flared disk
and weak outflow remain. While there is less envelope to attenuate the UV
radiation, there is also less material to irradiate and accretion luminosity
of older sources decrease, which results in less prominent emission from UV
tracers on large scales in Class I (Section 4.3). Another example of a
transition of molecular appearance between Class 0 and I is SO: this molecule
traces the outflow in Class 0 and the disk-envelope interface in Class I
primarily.

Table 3: Summary of evolution of chemical tracers Component |  Class 0 |  Class I
---|---|---
Envelope |  Cold, dense envelope results in cold tracers (DCO+, N2H+, N2D+) |  Envelope dissipates, extent of cold tracers is much smaller
Jet |  O-bearing molecules (CO, SiO, SO, H2CO) are present in high-velocity bullets |  Molecular jet disappears, seen only in atomic and ionised gas
Outflow |  Ice sputtering (CH3OH, HNCO) and grain destruction (SO, SiO) tracers are present |  Decreased outflow mass and less dense envelope result in no tracers of sputtering and grain destruction, only faint CO remains
Cavity walls |  Prominent signs of UV-irradiated cavity walls (CN, C2H, c-C3H2) |  No prominent signs of hydrocarbons in cavity walls, CN still present
Hot core |  COMs and simple tracers of ice sublimation and high-temperature chemistry (H2S, OCS) |  Small extent plus disk shadowing results in less complexity, except for outbursting sources
Disk |  Warm disk with COMs, dust obscuration |  Colder disk molecules, Keplerian rotation seen in H2CO, C18O, CN in the disk surface.

##  5 Conclusions

In this work, we have presented an overview of the molecular tracers of the
physical components in protostellar Class 0 and Class I systems. Table 2 and
Fig. 11 present an overview of the results, which are summarized below.

  * •

Protostellar envelopes are primarily traced with the dense gas tracers such as
C18O and CO, and snowline tracers such as DCO+, N2D+, H13CO+. For this
component, it is essential to use observations that are sensitive to large
scales of $\sim$ 5000 au.

  * •

Protostellar outflows are separated into different components. High-velocity
molecular jets are traced by O-bearing molecules CO, SiO, SO, and H2CO.
Entrained outflow material can be probed by low-velocity CO, SiO, SO, H2CO,
HCN as well as molecules released from ices by sputtering, such as CH3OH,
CH3CHO, HNCO, CH3CN, and CH3OCHO.

  * •

Outflow cavity walls are pronounced in UV irradiation tracers: C2H, c-C3H2,
and CN, and high-density tracers such as 13CS.

  * •

Hot cores, the inner warm envelopes, are probed predominantly by COMs and
simple molecules released from ice mantles and produced by high-temperature
chemistry such as SO, H2S, H2CS, OCS, H13CN, HCOOH, and HNCO.

  * •

Embedded disks, which are most clearly seen in Class I sources, can be seen in
C17O and H2CO while CN traces their upper layers. CO isotopologues can be used
to probe disks in Class 0 sources. However, the confusion with the envelope
emission requires a detailed analysis of the kinematics to confirm that a
Keplerian disk is present.

Hydrocarbons and COMs are found to co-exist routinely in protostellar sources:
the former are present in the UV-irradiated cavity walls and the latter in the
hot core and in the outflows. It is plausible that PDR chemistry for the
formation of hydrocarbons is sufficient to explain the presence of molecules
such as C2H and c-C3H2, and that no warm carbon-chain chemistry is required,
which relies on abundant CH4 in ices. Observations of more complex carbon
chains in cavity walls and envelopes are necessary to understand whether the
WCCC plays a role there. JWST will be able to probe the CH4 ice content at
spatial scales comparable to ALMA, which will show if there are any
differences in the ice composition between sources.

Throughout this study, the availability of both low-lying molecular lines to
probe the cold extended envelope and outflow, as well as high-excitation lines
to trace compact hot cores has been a key for the analysis. The occurrence of
the same molecules in different physical components provides an opportunity to
study the processes involved in their formation and excitation. For example,
ice mantle tracers sputtered from the grains in outflows give useful insight
on the ice composition, whereas comparison with abundances of thermally
released COMs in the inner region is a powerful tool to probe whether gas-
phase formation routes contribute to complex molecules formation.

JWST-MIRI will provide unprecedented resolution and sensitivity in mid-IR that
will allow probing the emission from hot gas in shocks as well as ice
absorption features. JWST observations can then be combined with our
understanding of the kinematics and spatial origin of molecules revealed by
ALMA.

Acknowledgements. The authors are grateful to the anonymous referee for
comments that improved the quality of the paper. This paper makes use of the
following ALMA data: ADS/JAO.ALMA# 2017.1.01174.S, ADS/JAO.ALMA#
2017.1.01350.S, ADS/JAO.ALMA# 2017.1.01371.S, ADS/JAO.ALMA# 2017.1.01413.S,
ADS/JAO.ALMA# 2013.1.00726.S, ADS/JAO.ALMA# 2016.1.00710.S. ALMA is a
partnership of ESO (representing its member states), NSF (USA) and NINS
(Japan), together with NRC (Canada), MOST and ASIAA (Taiwan), and KASI
(Republic of Korea), in co-operation with the Republic of Chile. The Joint
ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ. Astrochemistry in
Leiden is supported by the Netherlands Research School for Astronomy (NOVA).
Ł.T. acknowledges the ESO Fellowship Programme. M.L.vG. acknowledges support
from the Dutch Research Council (NWO) with project number NWO TOP-1
614.001.751. C.L.H.H. acknowledges the support of the NAOJ Fellowship and JSPS
KAKENHI grants 18K13586 and 20K14527. The National Radio Astronomy Observatory
is a facility of the National Science Foundation operated under cooperative
agreement by Associated Universities, Inc. This research has made use of
NASA’s Astrophysics Data System Bibliographic Services. This research made use
of NumPy (Harris et al. 2020); Astroquery (Ginsburg et al. 2019); Astropy, a
community-developed core Python package for Astronomy (Astropy Collaboration
et al. 2018, 2013); Pandas (McKinney 2010, 2011); and Matplotlib (Hunter
2007).

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##  Appendix A Tables

Table 4: Specifications of observations

Project ID |  Configuration |  $\lambda$ |  Res. |  MRSa |  Calibrationb |  Bandpass |  Phase |  Flux |  Targets
---|---|---|---|---|---|---|---|---|---
2017.1.1350.S |  Band 6 (C-4) |  1.3 mm |  0$\aas@@fstack{\prime\prime}$40 $\times$ 0$\aas@@fstack{\prime\prime}$29 |  5″ |  5.4.0-68 |  J0423-0120 |  J0336+3218 |  J0510+1800 |  SMM3, TMC1
2017.1.1350.S |  Band 6 (7m) |  1.3 mm |  6$\aas@@fstack{\prime\prime}$4 $\times$ 6$\aas@@fstack{\prime\prime}$1 |  25″ |  5.4.0-68 |  J0423-0120 |  J0336+3218 |  J0510+1800 |  SMM3, TMC1, IRAS4B, BHR71, Emb25, B1c
2017.1.1174.S |  Band 6 (C-4) |  1.3 mm |  0$\aas@@fstack{\prime\prime}$58 $\times$ 0$\aas@@fstack{\prime\prime}$39 |  6″ |  5.1.1 |  J1751+0939 |  J1830+0619 |  J1751+0939 |  B1c, S68N, SMM3
2017.1.1174.S |  Band 3 (C-2) |  3 mm |  2$\aas@@fstack{\prime\prime}$8 $\times$ 1$\aas@@fstack{\prime\prime}$8 |  16″ |  5.1.1 |  J0238+1636 |  J0336+3218 |  J0238+1636 |  B1c, S68N
2017.1.1371.S |  Band 5 (C-5) |  2 mm |  0$\aas@@fstack{\prime\prime}$6 $\times$ 0$\aas@@fstack{\prime\prime}$4 |  4″ |  5.1.1 |  J0237+2848 |  J0336+3218 |  J0237+2848 |  L1448, B1c, B5IRS1, HH211
2017.1.1413.S |  Band 6 (C-4) |  1.3 mm |  0$\aas@@fstack{\prime\prime}$42 $\times$ 0$\aas@@fstack{\prime\prime}$28 |  6″ |  5.4.0 |  J0510+1800 |  J0438+3004 |  J0510+1800 |  TMC1, 04302, L1527, TMC1A, L1489
2013.1.00726.S |  Band 6 (C-4, C-2) |  1.3 mm |  0$\aas@@fstack{\prime\prime}$45 $\times$ 0$\aas@@fstack{\prime\prime}$35 |  10″ |  4.2.2 |  J1733-1304 |  J1751+0939 |  Titan |  SMM1, S68N, Emb8N
2016.1.00710.S |  Band 3 (C-5) |  3 mm |  0$\aas@@fstack{\prime\prime}$5 $\times$ 0$\aas@@fstack{\prime\prime}$4 |  7″ |  4.7.38335 |  J1751+0939 |  J1838+0404 |  J1838+0404 |  SMM1, S68N, Emb8N

222 $a$$a$footnotetext: MRS - maximum recoverable scale at a given
configuration and wavelength $b$$b$footnotetext: Version of CASA used for
calibration

Table 5: Molecular transitions targeted Molecule |  Transition |  Frequency |  Eup |  $A_{ij}$ |  Serpens |  |  |  |  Perseus |  |  |  |  Taurus |  |  |  |
---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---|---
|  |  GHz |  K |  s-1 |  S68N |  SMM3 |  SMM1 |  8N |  B1c |  L1448 |  HH211 |  B5 |  04302 |  TMC1 |  TMC1A |  L1527 |  L1489
CO |  2–$1$ |  230.5380 |  17 |  6.9e-7 |  ✓ |  ✓ |  ✓ |  ✓ |  — |  — |  — |  — |  — |  ✓ |  — |  — |  —
C18O |  2–$1$ |  219.5604 |  16 |  6.1e-7 |  ✓ |  ✓ |  ✓ |  ✓ |  — |  — |  — |  — |  — |  ✓ |  — |  — |  —
13CO |  2–$1$ |  220.3987 |  16 |  6.1e-7 |  — |  ✓ |  — |  — |  — |  — |  — |  — |  ✓ |  — |  — |  — |
C17O |  21.5–11.5 |  224.7147 |  16 |  4.5e-7 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
H13CO+ |  1–0 |  86.7543 |  4 |  3.9e-5 |  ✗ |  — |  ✓ |  ✗ |  — |  — |  — |  — |  — |  — |  — |  — |  —
|  2–1 |  173.5067 |  13 |  3.7e-4 |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  — |  — |  — |  — |  —
|  3–2 |  260.2553 |  25 |  1.3e-3 |  ✓ |  — |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
H2CO |  30,3–20,2 |  218.2222 |  21 |  2.8e-4 |  ✓ |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  — |  —
|  31,2–21,1 |  225.6978 |  34 |  2.8e-4 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
|  32,1–22,0 |  218.7601 |  68 |  1.6e-4 |  — |  ✓ |  — |  — |  — |  — |  — |  — |  — |  ✓ |  — |  — |  —
|  91,8–91,9 |  216.5687 |  174 |  7.2e-6 |  — |  ✗ |  — |  — |  — |  — |  — |  — |  — |  ✗ |  — |  — |  —
H2CS |  51,4–41,3 |  174.3452 |  38 |  7.3e-5 |  — |  — |  — |  — |  ✓ |  ✓ |  ✗ |  ✓ |  — |  — |  — |  — |  —
|  70,7–60,6 |  240.2669 |  46 |  2.1e-4 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✗ |  ✗ |  ✗ |  ✓
SO |  56–45 |  219.9494 |  35 |  1.4e-4 |  — |  ✓ |  — |  — |  — |  — |  — |  — |  — |  ✓ |  — |  — |  —
|  67–56 |  261.8437 |  48 |  2.3e-4 |  ✓ |  — |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
SiO |  4–3 |  173.6883 |  21 |  2.6e-4 |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✗ |  — |  — |  — |  — |  —
|  5–4 |  217.1050 |  31 |  5.2e-4 |  ✓ |  ✓ |  ✓ |  ✓ |  — |  — |  — |  — |  — |  ✗ |  — |  — |  —
DCO+ |  3–2 |  216.1126 |  21 |  7.6e-4 |  ✓ |  ✓ |  ✓ |  ✓ |  — |  — |  — |  — |  — |  ✓ |  — |  — |  —
N2D+ |  3–2 |  231.3218 |  22 |  7.1e-4 |  — |  ✓ |  — |  — |  — |  — |  — |  — |  — |  ✗ |  — |  — |  —
OCS |  19–18 |  231.0610 |  111 |  3.6e-5 |  — |  ✓ |  — |  — |  — |  — |  — |  — |  — |  ✗ |  — |  — |  —
O13CS |  18–17 |  218.1990 |  99 |  3.0e-5 |  ✓ |  — |  ✓ |  ✗ |  — |  — |  — |  — |  — |  — |  — |  — |  —
H2S |  22–21 |  216.7104 |  83 |  4.8e-5 |  — |  ✓ |  — |  — |  — |  — |  — |  — |  — |  ✗ |  — |  — |  —
HNCO |  50,5–40,4 |  109.9058 |  16 |  1.8e-5 |  ✓ |  ✓ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  110,11–100,10 |  241.7741 |  70 |  2.0e-4 |  — |  — |  — |  — |  — |  — |  — |  — |  ✗ |  ✗ |  ✗ |  ✗ |  ✗
|  120,12–110,11 |  263.7487 |  82 |  2.6e-4 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  121,12–111,11 |  262.7696 |  125 |  2.6e-4 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
H2CCO |  131,13–121,12 |  260.1920 |  101 |  2.0e-4 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  132,11–122,10 |  262.7609 |  141 |  2.0e-4 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
HCOOH |  50,5–40,4 |  111.7468 |  16 |  1.5e-5 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  52,4–42,3 |  112.2871 |  29 |  1.3e-5 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  53,3–43,2 |  112.4596 |  45 |  9.7e-6 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  120,12–110,11 |  262.1036 |  83 |  2.0e-4 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
13CS |  5–4 |  231.2207 |  33 |  2.5e-4 |  ✓ |  ✓ |  ✓ |  ✓ |  — |  — |  — |  — |  — |  ✗ |  — |  — |  —
CN |  $1_{0,0.5,0.5}$–$0_{0,0.5,1.5}$ |  113.1442 |  5 |  1.1e-5 |  ✓ |  ✓ |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  $2_{0,1.5,2.5}$–$1_{0,1.5,2.5}$ |  226.3599 |  16 |  1.6e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✗ |  ✓ |  ✓ |  ✓ |  ✓
|  $2_{0,1.5,0.5}$–$1_{0,0.5,1.5}$ |  226.6166 |  16 |  1.1e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✗ |  ✓ |  ✗ |  ✗ |  ✓
|  $2_{0,1.5,2.5}$–$1_{0,0.5,1.5}$ |  226.6596 |  16 |  9.5e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
|  $2_{0,1.5,0.5}$–$1_{0,0.5,0.5}$ |  226.6637 |  16 |  8.5e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
|  $2_{0,1.5,1.5}$–$1_{0,0.5,0.5}$ |  226.6793 |  16 |  5.3e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
|  $2_{0,2.5,2.5}$–$1_{0,1.5,1.5}$ |  226.8742 |  16 |  9.6e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
|  $2_{0,2.5,3.5}$–$1_{0,1.5,2.5}$ |  226.8748 |  16 |  1.1e-4 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
|  $2_{0,2.5,1.5}$–$1_{0,1.5,0.5}$ |  226.8759 |  16 |  8.6e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
|  $2_{0,2.5,1.5}$–$1_{0,1.5,1.5}$ |  226.8874 |  16 |  2.7e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
|  $2_{0,2.5,2.5}$–$1_{0,1.5,2.5}$ |  226.8921 |  16 |  1.8e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✓ |  ✓
HCN |  1–0 |  88.6318 |  4 |  2.4e-5 |  ✓ |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  — |  —
H13CN |  1–0 |  86.3423 |  4 |  2.2e-5 |  ✓ |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  — |  —
|  2–1 |  172.6780 |  12 |  1.6e-4 |  — |  — |  — |  — |  ✓ |  ✓ |  ✓ |  ✗ |  — |  — |  — |  — |  —
CCH |  $3_{2.5,2}$–$2_{1.5,2}$ |  262.0788 |  25 |  6.5e-6 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  $3_{2.5,3}$–$2_{1.5,2}$ |  262.0648 |  25 |  5.3e-5 |  ✓ |  — |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  $3_{2.5,2}$–$2_{1.5,1}$ |  262.0673 |  25 |  4.8e-5 |  ✓ |  — |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  $3_{3.5,3}$–$2_{2.5,2}$ |  262.0064 |  25 |  5.5e-5 |  ✓ |  — |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  $3_{3.5,4}$–$2_{2.5,3}$ |  262.0042 |  25 |  5.7e-5 |  ✓ |  — |  — |  ✓ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
c-C3H2 |  61,6–50,5 |  217.8221 |  39 |  5.9e-4 |  — |  ✓ |  — |  — |  — |  — |  — |  — |  — |  ✓ |  — |  — |  —
|  72,6–71,7 |  218.7327 |  61 |  9.8e-5 |  — |  ✓ |  — |  — |  — |  — |  — |  — |  — |  ✓ |  — |  — |  —
|  74,3–73,4 |  112.4908 |  83 |  4.5e-5 |  ✓ |  ✓ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
CH3CHO |  61,6,1–51,5,1 |  112.2487 |  21 |  4.5e-5 |  ✓ |  ✓ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  61,6,0–51,5,0 |  112.2545 |  21 |  4.5e-5 |  ✓ |  ✓ |  — |  ✗ |  — |  — |  — |  — |  — |  — |  — |  — |  —
|  94,5–84,4 |  173.5191 |  78 |  1.4e-4 |  — |  — |  — |  — |  ✓ |  ✓ |  ✗ |  ✗ |  — |  — |  — |  — |  —
|  140,14–130,13 |  262.9601 |  96 |  6.2e-4 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
CH3CN |  60–50 |  110.3835 |  19 |  1.1e-4 |  ✓ |  ✓ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  61–51 |  110.3814 |  26 |  1.1e-4 |  ✓ |  ✓ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  62–52 |  110.3750 |  47 |  9.9e-5 |  ✓ |  ✓ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  63–53 |  110.3644 |  83 |  8.3e-5 |  ✓ |  ✓ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  64–54 |  110.3495 |  133 |  6.2e-5 |  ✓ |  ✗ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
CH3OH |  21,0–10,1 |  261.8057 |  28 |  5.6e-5 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  50–40 |  241.7914 |  35 |  6.1e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✓ |  ✗ |  ✗ |  ✗ |  ✓
|  5-1–4-1 |  241.7672 |  40 |  5.8e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✗ |  ✗ |  ✗ |  ✗ |  ✓
|  52–42 |  241.9046 |  57 |  5.0e-5 |  — |  — |  — |  — |  — |  — |  — |  — |  ✗ |  ✗ |  ✗ |  ✗ |  ✓
|  42–51 |  234.6834 |  61 |  1.9e-5 |  — |  ✗ |  — |  — |  — |  — |  — |  — |  — |  ✗ |  — |  — |  —
13CH3OH |  21,1–10,1 |  259.9865 |  28 |  5.5e-5 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  52,3–41,3 |  261.8057 |  56 |  7.4e-5 |  ✓ |  — |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
CH3OCHO |  95,5–85,4 |  110.8805 |  43 |  1.4e-5 |  ✓ |  ✗ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  100,10–90,9 |  111.1699 |  30 |  2.0e-5 |  ✓ |  ✗ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
|  91,8–81,7 |  111.6741 |  28 |  2.0e-5 |  ✓ |  ✗ |  — |  ✗ |  ✓ |  — |  — |  — |  — |  — |  — |  — |  —
333✓- transition detected, ✗- transition not detected, – transition not
targeted

##  Appendix B Continuum images

##  Appendix C Envelope tracers

##  Appendix D Outflow

##  Appendix E Cavity walls

##  Appendix F Inner envelope

##  Appendix G Additional plots

f