# Investigating Anomalous Photochemistry in the Inner Wind of IRC+10216
Through Interferometric Observations of HC3N

Mark A. Siebert  Department of Astronomy, University of Virginia,
Charlottesville, VA 22904, USA  Marie Van de Sande  School of Physics and
Astronomy, University of Leeds, Leeds, UK  Thomas J. Millar  School of
Mathematics and Physics, Queen’s University Belfast, Belfast, UK  Anthony J.
Remijan  National Radio Astronomy Observatory, Charlottesville, VA 22903, USA

(Accepted October 26, 2022)

###### Abstract

In recent years, many questions have arisen regarding the chemistry of
photochemical products in the carbon-rich winds of evolved stars. To address
them, it is imperative to constrain the distributions of such species through
high angular resolution interferometric observations covering multiple
rotational transitions. We used archival ALMA observations to map rotational
lines involving high energy levels of cyanoacetylene (HC3N) toward the inner
envelope (radius $<8"$/$1000$ AU) of the carbon star IRC+10216. The observed
lines include the $J=28-27$, $J=30-29$, and $J=38-37$, transitions of HC3N in
its ground vibrational state. In contrast to previous observations of linear
carbon chains toward this AGB star which show extended, hollow emission at
$15"-20"$ radii (e.g. C4H, C6H, HC5N), the maps of the HC3N lines here show
compact morphologies comprising various arcs and density enhancements, with
significant emission from gas clumps at an angular distance of ${\sim}3”$ (350
AU) from the central AGB star. We compared visibility sampled non-LTE
radiative transfer models with the observed brightness distributions, and
derive a fractional abundance with respect to H2 of $10^{-8}$ for HC3N at the
radii probed by these lines. These results are consistent with enhanced
photochemistry occurring in warm (${\sim}$200 K) regions of the circumstellar
envelope. After application of a specialized chemical model for IRC+10216, we
find evidence that the enhanced HC3N abundances in the inner wind are most
likely due to a solar-type binary companion initiating photochemistry in this
region.

Astrochemistry – circumstellar matter – line: identification – stars: AGB,
individual (IRC+10216)

††journal: ApJ

##  1 Introduction

IRC+10216, also known as CW Leo, is a well-studied nearby ($d\sim 130$ pc,
Men’shchikov et al., 2001) Asymptotic Giant Branch (AGB) star with a
chemically rich expanding circumstellar envelope (CSE). The CSE is the result
of dust-driven mass loss at a rate of ${\sim}2\times 10^{-5}$ M☉ yr-1(De Beck
et al., 2012), and represents one of the final stages of evolution for 1–8 M☉
stars immediately preceding the transition between the AGB and the formation
of a planetary nebula (PN). During this stage, depending on how much carbon
has been convectively dredged up from the star’s core relative to oxygen
(Weiss & Ferguson, 2009), AGB stars can either be M-type (C/O$<1$), C-type
(C/O$>1$), or S-type (C/O${\sim}1$). IRC+10216 has a C/O ratio of 1.4 (Winters
et al., 1994), which allows for a complex system of carbonaceous gas phase and
dust grain chemistry to occur throughout its extended envelope. To date, more
than 95 unique molecules have been identified toward IRC+10216 (McGuire,
2022). Due to its close proximity and history as a molecular line object,
IRC+10216 forms the basis of our understanding on how carbon-rich stellar
material is processed and returned to the interstellar medium by low-to-
intermediate mass stars.

An important aspect of AGB evolution that has become clear with the advent of
ALMA is the prevalence of binary companion stars influencing mass loss and
shaping circumstellar winds (Decin, 2021). Companions are difficult to detect
directly in these systems, due to the intrinsic luminosity of the AGB star,
and the high attenuation of visible light by their dusty envelopes. However,
several recent studies of gas and dust kinematics have shown the perturbative
effects binary companions can have on many AGB and proto-planetary nebula
(PPNe) systems. Both small-scale asymmetries, such as density-enhanced clumps
(e.g. Khouri et al., 2016; Agúndez et al., 2017; Leão et al., 2006), and
large-scale structures, such as spirals (e.g. Mauron & Huggins, 2006; Maercker
et al., 2016; Decin et al., 2015; Ramstedt et al., 2017), disk geometries (e.g
Kervella et al., 2014; Homan et al., 2018; Sahai et al., 2022; Homan et al.,
2020), and bipolar outflows (e.g. Kim et al., 2015b; Sahai et al., 2017;
Lagadec et al., 2011) have been detected. In particular, the ALMA Large
Program ATOMIUM (2018.1.000659.L, PI: Decin) has shown the widespread nature
of these structures through a survey of numerous AGB outflows (Decin et al.,
2020). It is now thought that influence from stellar and planetary companions
is closely related to the stark morphological changes in circumstellar
material during the transition between the AGB and PN stages of evolution
(Sahai et al., 2007). IRC+10216 shows signs of binary influence, based on the
interspaced shells in its extended envelope (Cernicharo et al., 2015a; Guélin
et al., 2018), spiral structure in its inner layers (Decin et al., 2015), as
well as the detection of a tentative compact source $0.5$″ from the star via
archival HST monitoring by Kim et al. (2015a, 2021).

While their physical effects have been studied in detail, the chemical impact
of companion stars embedded in CSEs and PPNe has not yet been observationally
tested. The typical framework for how gas phase chemistry proceeds in
circumstellar envelopes assumes precursor molecules (e.g. CO, HCN, SiO) form
out of a combination of equilibrium and shock-induced chemistry in the high
temperature regions near the stellar photosphere (Agúndez et al., 2020;
Cherchneff, 2012). Once they reach the outer regions of the envelope where the
density of attenuating dust grains is lower, UV photons from the interstellar
radiation field (ISRF) initiate a kinetically-driven photochemistry to form a
variety of exotic product molecules (e.g. HC7N, SiN, MgNC) (Ziurys, 2006; Li
et al., 2016; Millar, 2020). While this scheme is consistent with the measured
abundances of many molecules, recent studies of IRC+10216 have revealed that
this picture may be much more complicated than previously thought. Agúndez et
al. (2015) and Quintana-Lacaci et al. (2017) discovered that CH3CN and NaCN
are present in the inner envelope ($<5$″). This was an unexpected result as
the proposed formation routes of both these molecules require photochemistry,
which typically takes hold further out (10″–25″) in the CSE (Dinh-V-Trung &
Lim, 2008). Similar results were also found for the product molecules water
(H2O) and diacetylene (C4H2) (Agúndez et al., 2010; Fonfría et al., 2018)

One molecule that is central to photochemistry in CSE environments is
cyanoacetylene (HC3N). In the gas phase, this species is formed primarily
through the reaction between acetylene (C2H2) and the cyanide radical (CN),
the latter of which is only available as a reactant when the flux of UV
photons is large enough to efficiently dissociate the precursor molecule HCN
(Agúndez et al., 2017). Because C2H2 is an abundant precursor molecule and one
of the main carriers of carbon in C-rich CSEs (Santoro et al., 2020), HC3N is
one of the first product molecules to form out of UV-driven chemistry, making
it useful in probing the photoprocessing of gas phase material in these
objects.

HC3N and longer cyanopolyynes in the envelope of IRC+10216 were studied with
the VLA by Dinh-V-Trung & Lim (2008) and subsequently with ALMA by Agúndez et
al. (2017), where it was seen between 10″–20″(1300–2600 AU). This is where the
HC3N abundance is expected to peak based on photochemical models of this CSE;
however, these observations targeted low-lying rotational transitions
(including states up to $J=12$ and $E_{\mathrm{up}}=34$ K) which naturally
trace lower temperature regions. In order to gain a more complete view of this
molecule in the envelope of IRC+10216, and to better understand the onset of
photochemistry in these environments, it is important to spatially constrain
higher energy transitions of HC3N that could be discernable in the warmer
regions where CH3CN and NaCN were observed.

To this end, we present an archival study of IRC+10216 with the goal of
characterizing HC3N in the inner (arcsecond scale) layers of its envelope. We
utilize observations from three separate ALMA projects to analyze the spectra
and spatial distributions of the $J=28-27$, $J=30-29$, and $J=38-37$ lines of
HC3N toward IRC+10216. Using maps of these lines at the systemic velocity
combined with radiative transfer models, we quantify the abundance of
cyanoacetylene in the inner regions of IRC+10216 and discuss how they affect
our understanding of photochemical processing in this CSE. In Section 2 we
describe the observations and reduction procedure, in Section 3.2 we present
emission maps, in Section 3.3 we use non-LTE radiative transfer models to
derive the abundance profile of HC3N, and in Section 4 we utilize a
specialized photochemistry model to compare our observations with kinetic
theory and test the hypothesis that molecules surrounding IRC+10216 are
affected by a binary stellar companion embedded in the inner envelope.

##  2 Observations and reduction

Table 1: Summary and characteristics of the ALMA data utilized in this work.

ALMA Project Code | Obs. Date(s) | Array(s) | Synthesized beam |  |  HC3N line | Rest Frequency | $E_{\mathrm{up}}$ | Image RMS
---|---|---|---|---|---|---|---|---
|  |  | ($\theta_{\mathrm{maj}}\times\theta_{\mathrm{min}}$) |  | ($J^{\prime}\rightarrow J^{\prime\prime}$) | (MHz) | (K) | (mJy/beam)
2019.1.00507.S | Nov. 2019 – Mar. 2020 | 12m+7m | $0.480"\times 0.449"$ |  | $28\rightarrow 27$ | 254 699.5 | 177 | 1.8
2011.0.00229.S | Apr. 2012 | 12m | $0.725"\times 0.520"$ |  | $30\rightarrow 29$ | 272 884.7 | 203 | 7.2
2016.1.00251.S | May 2018 | 12m | $0.848"\times 0.761"$ |  | $38\rightarrow 37$ | 345 609.0 | 324 | 1.9

Note. — Observations from ALMA Project 2019.1.00507.S include two 12m
configurations and two from the 7m ACA. Transition frequencies and upper state
energies were obtained from Cologne Database of Molecular Spectroscopy (CDMS)
(Müller et al., 2005). Reported RMS values were measured from an emission-free
channel of the final reconstructed image cube.

Figure 1: Coverage of the uv-plane offered by the three ALMA projects utilized
in this work. Baselines from compact 12m configurations are shown in blue,
those from extended 12m configurations are shown in brown, and 7m ACA
baselines are shown in orange. Observations from left to right were used to
image the $J=28-27$, $J=30-29$, and $J=38-37$ transitions of HC3N,
respectively.

We utilize observations from three publicly available archival ALMA projects
to map warm HC3N emission toward IRC+10216. Fig. 1 demonstrates the different
antenna configurations and coverage of the uv-plane for the data sets, while
Table 1 summarizes these different observations and the target HC3N
transitions for this work. Project 2019.1.00507.S provides several spectral
scans of IRC+10216 at Band 6 between 250 and 270 GHz using both the main 12m
array and the 7m Atacama Compact Array (ACA). These observations were used by
He et al. (2019) to study the time-variability of millimeter lines toward
IRC+10216, and cover the $J=28-27$ transition of HC3N. In order to recover
flux from as many spatial scales as possible, we combined visibility
measurements from four separate array configurations included in this project,
including two 12m executions with baseline lengths ranging from 15 m to 1230 m
and two ACA measurements covering baselines between 9 and 50 m (see Fig. 1).
Data combination was performed with antenna-specific weighting, and the
resulting images have a maximum recoverable scale of $30$″.

Project 2011.0.00229.S provides a Cycle 0 spectral line survey of IRC+10216
covering most of Band 6 using the 12m array, allowing emission from the
$J=30-29$ transition to be mapped as well; however, only 18 antennas were
available at this time, as opposed to the 42 that were used for the later data
sets. These data cover baselines ranging from 27 m to 402 m with a maximum
recoverable scale of ${\sim}1"$.

Figure 2: Spectra demonstrating the three HC3N lines studied in this work
(red), in addition to nearby and interloping molecular lines (gray). Spectra
were created from image cubes using a 30″aperture centered on the position of
IRC+10216. The lower x-axes denote sky frequencies, while the upper axis shows
velocity using the rest frequency of the respective HC3N transition. All
transitions are from the ground vibrational state unless indicated otherwise.

Project 2016.1.00251.S provides a Band 7 observation between 330 and 346 GHz
using the 12m array. This ALMA configuration includes baselines between 15 m
and 314 m, and offers a maximum recoverable scale of $8$″. These data were
used to map the $J=38-37$ line of HC3N in addition to the $J=3-2$ line of
^13CO at $\nu_{\mathrm{rest}}=330\,588.0$ MHz. The ^13CO observations are
useful in qualitatively characterizing the overall envelope density structure.

As Figure 1 demonstrates, the ALMA Projects we utilize have very different
antenna setups, and consequently disparate spatial frequencies sampled. Due to
these differences between instrument configurations and the non-uniqueness of
the image reconstruction process, care must be taken when directly comparing
brightness distributions between data sets. All observations were pointed near
the continuum peak of IRC+10216 at $\alpha_{\mathrm{J2000}}$ =
09$\fh$47$\fm$57.458$\fs$, $\delta_{\mathrm{J2000}}$ = +13°16′43.90″, despite
having slightly different phase centers between projects. Errors in pointing
are estimated to be 30 mas for the $J=28-27$ observations, and 40 mas for the
$J=30-29$ and $J=38-37$ observations111Referencing Section 10.5.2 of the ALMA
Technical Handbook (Cortes et al., 2020).

All data were calibrated using the standard calibration pipeline of the Common
Astronomy Software Application (CASA) (McMullin et al., 2007). Visibilities
from the Band 6 data sets were resampled to a spectral resolution of 1.7 km
s-1 to match the correlator configuration in the Band 7 project. Continuum
subtraction was done in visibility space using line free channels present in
each spectral window containing HC3N. Then, imaging was performed with the
TCLEAN casatask, using Briggs weighting with a 0.5 robust parameter for the
$J=28-27$ and $J=38-37$ lines, and natural weighting for the $J=30-29$ line
due to its lower signal-to-noise. Iterative masking was done with user-defined
regions to optimize the deconvolution process. In the case of the $J=28-27$
line, multi-scale cleaning was applied to account for the multiple combined
array configurations (Cornwell, 2008). All rest frequencies were obtained from
Cologne Database of Molecular Spectroscopy (CDMS) (Müller et al., 2005), and
shifted to the systemic velocity of IRC+10216, $V_{sys}=-26.5$ km
s-1(Cernicharo et al., 2000).

##  3 Results

###  3.1 Spectra, interloping lines, and flux recovery

Integrated spectra from each synthesized image cube were obtained using a
30″circular aperture centered on IRC+10216. These spectra are shown in Figures
2a-c. In these, we see emission in the characteristic “double-horned” line
profile centered on the systemic velocity of $-26.5$ km s-1, as well as some
contributions from interloping lines. For the $J=28-27$ line of HC3N, a very
close blend is seen with the $J(K_{a},K_{c})=4(0,4)-3(0,3)$ transition of
methanimine (CH2NH) at $\nu_{\mathrm{rest}}=254\,685.1$ MHz. This line appears
14 MHz (17 km s-1) from the central frequency of HC3N. In the spectrum (Fig.
2a.), a double horn profile of CH2NH can be seen with the blue-shifted peak
falling very close to the central channel of HC3N. This blend was also
detected by Tenenbaum et al. (2010). Since the aim of this study is to
characterize the radial abundance using channel maps made at the systemic
velocity, we must take care to make sure we do not attribute blue-shifted
CH2NH emission to HC3N. However, considering the geometry of spherical
outflows and empirical evidence from previous ALMA studies of AGB stars (e.g.
channel maps in Decin et al. (2018) and Brunner et al. (2018)), we expect that
the contamination from the highest velocity emission is limited only to the
smallest spatial scales and appears very close to the position of IRC+10216.

In the case of the $J=38-37$ line of HC3N (Fig. 2c.), there is a slight blend
with the $J(K_{a},K_{c})=15(14,1)-15(14,0)$ transition of SiC2 at
$\nu_{\mathrm{rest}}=345\,575.2$ MHz. This lies 34 MHz (29 km s-1) away from
the central frequency of our target HC3N line, so with a half line width of
$14.5$ km s-1, only the most red-shifted channels of the image cube are
contaminated by emission from SiC2.

The $J=28-27$ and $J=30-29$ lines of HC3N were also observed in the single
dish line surveys of He et al. (2008) and Tenenbaum et al. (2010), both using
the Arizona Radio Observatory 10 m Submillimeter Telescope (SMT). These works
provide useful data points that allow us to characterize the amount of spatial
filtering in our interferometric measurements of the same lines. Both He et
al. (2008) and Tenenbaum et al. (2010) detected the $J=28-27$ line of HC3N
with a peak brightness temperature $T_{R}=130$ mK. Considering that the SMT
has a HPBW of 30″at 254.7 GHz, we can compare this directly with the spectrum
in Fig. 2a., which peaks at 120 mK. In contrast, the $J=30-29$ line shown in
Fig. 2b. was detected in Tenenbaum et al. (2010) with a peak temperature of
$T_{R}=90$ mK, while we instead obtain a value of 10 mK. It is clear that the
observations from Project 2019.1.00507.S, which comprise both ACA and multiple
12m executions spanning a large range of baselines, are able to recover more
than $90\%$ of the HC3N ($J=28-27$) emission. However the Cycle 0 observations
of the $J=30-29$ line exhibit a large amount of spatial filtering due to the
fewer number of antennas and poorer uv-coverage. Unfortunately, there exist no
published single-dish measurements for the $J=38-37$ transition of HC3N toward
IRC+10216, but we still take steps to account for the flux loss of these
measurements described in Section 3.3.

###  3.2 Emission Maps

Figure 3: ALMA emission maps of three HC3N lines toward IRC+10216. Images were
created at the systemic velocity of -26.5 km s-1with a channel width of 1.69
km s-1. IRC+10216 is located at the center of the maps at
$\alpha_{\mathrm{J2000}}$ = 09$\fh$47$\fm$57.458$\fs$,
$\delta_{\mathrm{J2000}}$ = +13°16′43.90″, and a reference circle with
diameter 15″is shown in yellow. Because data from separate observation sets
are used, the beam sizes and absolute brightness scales are different for each
panel. The shapes of the synthesized beams are shown in the bottom left
corner, with sizes 0.48″$\times$0.45″, 0.73″$\times$0.52″, and
0.85″$\times$0.76″from top to bottom.

Reconstructed maps of the three studied HC3N lines are shown in Figure 3. All
images were made at the systemic velocity of IRC+10216 ($V_{sys}=-26.5$ km
s-1). These maps show morphologies centered on the position of the AGB star
with emission primarily within a radius of $5$″. The distribution of the
$J=28-27$ line includes an extended component out to 3″surrounded by several
asymmetric arcs and clumps (Fig. 3a.). Due to the higher temperatures probed
by this transition, the majority of its emission traces regions entirely
separated from that observed in the lower excitation lines mapped by Agúndez
et al. (2017) and Dinh-V-Trung & Lim (2008) ($J=10-9$ and $J=5-4$; which peak
at radius $14"$), with the notable exception of a region to the SW at a radius
of $10"$. In the central component (within $3"$), the average brightness of
the $J=28-27$ line is 10 mJy beam-1, with notable bright clumps to the N and
SW. Although a central point source is seen at the position of the carbon
star, we expect that this is contributed from the most blue-shifted channel of
the weak interloping line of CH2NH (see Section 3.1 and Fig. 2), since this
peak is notably absent from the unblended image in Fig. 3c.

Figure 4: Contour maps of HC3N $J=28-27$ $J=30-29$, and $J=38-37$ (white)
overlaid on the image of ^13CO $J=3-2$ toward IRC+10216. All components are
again taken at the systemic velocity and centered at the position of IRC+10216
$\alpha_{\mathrm{J2000}}$ = 09$\fh$47$\fm$57.458$\fs$,
$\delta_{\mathrm{J2000}}$ = +13°16′43.90″, and the yellow circle denotes a
radius of 7.5″. Contour levels for each line are shown at 3, 6, 9, and 12
times the rms.

The $J=38-37$ line (Fig. 3c.) also shows bright, spatially resolved emission
that is entirely distinct from the maps of HC3N in Agúndez et al. (2017). In
contrast to $J=28-27$, the distribution of this line is entirely concentrated
within a radius of $4"$ and shows no central peak on the position of
IRC+10216, since there is no contaminating emission present in this channel.
The imaged line is again not azimuthally symmetric, and traces the same bright
clumps and arcs seen in the $J=28-27$ line. It reaches a peak brightness of
$122\pm 1.9$ mJy beam-1 in the SW clump at $2"$. The more compact morphology
can be qualitatively explained by the higher upper state energy (323 K) and
larger (by a factor of 3) Einstein $A_{ij}$ coefficient of this transition
relative to the $J=28-27$ line.

Finally, the Fig. 3b. shows the $J=30-29$ line, which also appears toward the
same bright clump ${\sim}2$″to the SW. However due to the sensitivity of these
data, very little additional emission is seen from this transition. If this
line has a similar extended component as seen in $J=28-27$, we would expect
most of this flux to be resolved out by the interferometer given its maximum
recoverable scale. However, if the $J=30-29$ line had a compact, point-like,
component at the position of IRC+10216 with a similar peak brightness to the
SW clump (as is the case with $J=28-27$), we would expect to detect that even
at the limited sensitivity of these data. Due to the similar excitation
conditions of the $J=30-29$ and $J=28-27$ lines, we take this as further
evidence of the central peak in Fig. 3a) being contributed by CH2NH, and not
HC3N. Nevertheless, because of the observational limitations of this ALMA
project (2011.0.00229.S), we do not include this line in our radiative
transfer analysis.

Figure 4 shows the observed HC3N lines as contours overlaid on the image of
^13CO $J=3-2$ at $\nu_{\mathrm{rest}}=330\,587.97$ MHz created from the Band 7
ALMA observations. Using a CO abundance of $6\times 10^{-4}$ (De Beck et al.,
2012; Agúndez et al., 2012a) and a ^12C/^13C ratio of 48 (Knapp & Chang,
1985), we adopt a constant ^13CO abundance of $1.3\times 10^{-5}$ with respect
to H2 throughout the regions of IRC+10216 we are investigating. Therefore, it
is a useful indicator of the density structure. Keeping in mind the spatial
filtering effects due to the lack of very short antenna spacings in these
data, we see that the structures traced by HC3N are spatially coincident with
density enhancements resulting from non-isotropic mass-loss. The clump SW of
IRC+10216 which is seen in all three observed HC3N lines is also apparent here
as the brightest area of ^13CO emission away from the central point source
(see Fig. 4). Between the bright arcs we see several areas of negative flux,
which are image artifacts resulting from resolved out emission and not
indicative of actual absorption. When compared with images in Agúndez et al.
(2015), we find that both HC3N and ^13CO trace similar compact structures to
CH3CN ($J_{K}=14_{3}-13_{3})$, including the SW clump.

###  3.3 Radiative Transfer Modeling

In order to draw conclusions regarding the abundance of HC3N, we must first
consider the excitation physics of this molecule in the inner envelope of
IRC+10216. In many cases, Local Thermodynamic Equilibrium (LTE) is a good
approximation for chemical studies of CSEs due to their high central
densities. However, the rotational transitions observed here have
comparatively large Einstein $A_{ij}$ coefficients ($10^{-2.9}-10^{-2.5}$
s-1), so it is possible that they are sub-thermally excited even at arcsecond
scales. In Figure 5, we show the predicted excitation temperatures of the
three HC3N transitions as a function of gas density for two different kinetic
temperatures, generated using the RADEX radiative transfer software (van der
Tak et al., 2007). Using the 1D density law from (Eq. 1 in Agúndez et al.
2012a) along with a mass-loss rate of $2.7\times 10^{-5}$ M☉ yr-1 (Guélin et
al., 2018), we also show the expected average gas densities at the radii where
most of the emission is seen (between 1″and 5″). From this, it is clear that
LTE could only be assumed if densities were higher than $10^{8}$ cm-3, and the
observed gas density is below the critical densities of all three HC3N
transitions in the mapped regions. Futhermore, pumping to vibrational states
via thermal infrared (IR) photons from dust grains has been shown to be an
important excitation mechanism in the envelope of IRC+10216 (Deguchi &
Uyemura, 1984; Cernicharo et al., 2014; Agúndez et al., 2017; Keady & Ridgway,
1993; Fonfría et al., 2008), so non-LTE calculations are necessary to take all
of these factors into account.

Figure 5: RADEX models of the three observed transitions of HC3N run with for
a range of densities at $T_{k}=30$ K (dashed lines) and $T_{k}=100$ K (solid
lines). The shaded region denotes the average gas densities at radii where we
observe these lines toward IRC+10216 (between 1″and 5″).

Therefore, we adopt a forward modeling approach in predicting the expected
brightness distributions and determining the abundance of HC3N needed to
reproduce the maps in Fig. 3. We utilize the LIne Modeling Engine (LIME) to
simulate the level populations of HC3N, solve the radiative transfer equation,
and produce synthetic image cubes (Brinch & Hogerheijde, 2010). LIME uses
density-weighted grid points distributed in a computational space using
Delaunay triangulation. Populations are solved based on the local conditions
of the corresponding Voronoi cells, as well as the incident radiation from
adjacent cells. For all our calculations, we use 50 000 grid points in a
computational volume with radius 130 000 AU (1000″). Though LIME performs
these calculations in 3D space, we use a 1D approximation for the physical
structure of IRC+10216, since the 3D asymmetries apparent in Figures 3 and 4
are not presently constrained in their local physical conditions and require
hydrodynamical simulations that are outside the scope of this work. In
general, we adopt the same model for the CSE of IRC+10216 as Agúndez et al.
(2012a). This includes the stellar parameters, velocity structure, gas and
dust temperature profiles, and density gradient. The only difference in our
physical model is the use of a slightly larger mass-loss rate of $2.7\times
10^{-5}$ M☉ yr-1 , as this value was measured by Guélin et al. (2018)
specifically using CO emission at radii within 15″, meaning it is better
suited for this study of warm HC3N emission. The gas density in the extended
envelope is given by

$n(r)=\frac{\dot{M}}{4\pi r^{2}v_{\infty}m_{g}}$ (1)

where $\dot{M}$ is the mass-loss rate, $r$ is the radius, $v_{\infty}$ is the
expansion velocity, and $m_{g}$ is the average gas particle mass taken as 2.3
amu considering H2, He, andCO (Agúndez et al., 2012a). So a change in
$\dot{M}$ has a significant effect on molecular volume densities and
corresponding line intensities. This should be kept in mind when comparing our
results with previous chemical studies of IRC+10216 that use the more typical
value of $2.0\times 10^{-5}$ M☉ yr-1 (Agúndez et al., 2017, 2015; Quintana-
Lacaci et al., 2017).

For spectroscopic data, the Leiden Atomic and Molecular Database (LAMDA)
provides collisional rates of ^13CO and HC3N with ortho- and para- H2 (van der
Tak et al., 2020). The ortho/para H2 ratio is assumed to be 3 at all locations
in the envelope. For HC3N, the collisional rates measured by Faure et al.
(2016) only go up to the $J=37$ rotational state, so we extrapolated these to
the $J=45$ level and use line strengths from CDMS for the added radiative
transitions (Müller et al., 2005). The $J=45$ level was chosen to account for
population in rotational states above $J=38$ while keeping the computational
load low. We note that the choice of this maximum level had little effect on
the predictions, as we found that running the model with maximum $J=40$ and
$J=50$ yielded nearly identical results.

In addition, we included the first excited states of the $\nu_{5}$ and
$\nu_{6}$ vibrational modes at 663 cm-1 and 499 cm-1 (Jolly et al., 2007)
using radiative rates from the HITRAN database (Gordon et al., 2022). The
effect of including these states is similar to the result of Massalkhi et al.
(2019), where emission is shifted to larger radii as cooler molecules could be
pumped to higher $J$ values by first being radiatively excited into these
vibrational states then subsequently de-exciting through the P-branch ($\Delta
J=+1$) (Costagliola & Aalto, 2010). Collisional transitions within vibrational
states were not considered, as the time scales of vibrational de-excitation
(${\sim}10^{-4}$ s) are several orders of magnitude shorter than the expected
collisional excitation rates (${\sim}10^{4}$ s). To model thermal emission
from dust grains, we use the opacity law of Ossenkopf & Henning (1994).

After obtaining ray-traced simulated image cubes from LIME, we test two
separate methods of post-processing before comparing with observations. The
first involves a simple image plane continuum subtraction followed by
convolution of the channel maps with the synthesized beams in Fig. 3. For the
second method, we utilize the Python module vis_sample222vis_sample is
publicly available under the MIT license at
https://github.com/AstroChem/vis_sample and described in further detail in
(Loomis et al., 2018) to simulate the incomplete sampling of sky brightnesses
using input uv grid points directly from the observation sets. The results are
simulated visibilities that were then continuum subtracted and imaged in the
same manner as their respective observations (described in Section 2). The
estimated flux error in the simulated image is 5%, based on the converged
signal-to-noise of the individual grid points. Because the input model is
spherically symmetric, we take the azimuthal average of the LIME-generated
channel map at the systemic velocity of IRC+10216 to compare the modeled
emission with ALMA images.

The result of both methods when applied to the $J=3-2$ line of ^13CO, for
which the abundance is known a priori, is shown in Fig. 6a. The convolved LIME
model of this line overpredicts the brightness, especially in the extended
layers of IRC+10216. In contrast, the visibility sampled and reconstructed
image reproduces the peak brightness to within 5%, and predicts the much lower
average brightness of this line due to spatial filtering. From this it is
clear that the latter method better accounts for the lack of data from short
baselines and the resulting flux loss at large spatial scales, so we employ
the visibility sampling routine for the case of HC3N as well.

####  3.3.1 Results for HC3N

Figure 6: Observed (solid lines) and LIME simulated (dashed lines) brightness
profiles of ^13CO, HC3N $J=38-37$, and HC3N $J=28-27$ from top to bottom.
Profiles were obtained by taking the azimuthal average of the channel map at
the systemic velocity of IRC+10216 centered on the position of the AGB star.
Shaded regions denote the standard deviation of the surface brightness in each
radial bin. In the top panel, the profile is also shown when a convolution
with the synthesized beam of observations is used instead of visibility
sampling. In the bottom panel, the expected contaminating emission from CH2NH
is shown as well (dotted line).

We use the above forward-modeling routine to simulate the $J=28-27$ and
$J=38-37$ lines of HC3N, now varying the abundance to reproduce the observed
radial brightness distribution. Primarily, we modify it at radii less than 4″,
since this is where we see the bulk of emission (shaded regions in Fig. 7);
however, we also adjust the peak abundance to reflect the observed lack of
emission at radii larger than 8″. We start with a function similar to the
modeled result in Agúndez et al. (2017), with a peak abundance of 1$\times
10^{-6}$ at $r=3\times 10^{16}$ cm that steadily decreases inwards, and drops
off steeply outwards. We then adjust the magnitude shape to match the
distribution and intensity, by eye, of azimuthally averaged emission for both
HC3N lines.

The resulting simulated intensity profiles are shown in Figures 6b. and 6c.
The corresponding best-fit abundance profile is shown in Fig. 7a., and the
volume density is shown in Fig. 7b. The resulting modeled intensity curves are
in good agreement with the observed azimuthal averages at the radii where we
see the bulk of HC3N emission. There is a slight disagreement in the regions
between 8″–14″where the average predicted $J=28-27$ emission increases 2 mJy
beam-1; however, this is very close to the rms of the observations, so it is
not expected that we would detect consistent emission at this level. We do
detect some clumps with a peak abundance of 25 mJy beam-1 in these regions, as
seen outside the yellow reference circle to the SW and NW in Fig. 3a., but no
extended component is found.

In the inner layers where the HC3N emission is primarily seen, an average
abundance of ${\sim}10^{-8}$ is needed, and it must be decreasing inward to
avoid centrally peaked channel maps. In typical fashion for characterizing the
radial abundance distributions of molecules forming in outer regions of CSEs
(e.g. Agúndez et al., 2015), we parameterize this with an offset Gaussian in
log radius:

$f\left(r\right)=f\left(r_{p}\right)\mathrm{exp}\left(-\left(\frac{\log r-\log
r_{p}}{\Delta\log r}\right)^{2}\right)$ (2)

where $f$ is the fractional abundance with respect to H2, $r_{p}$ is the peak
radius, and $\Delta\log r$ adjusts the width of the feature. For the retrieved
profile, we obtain $f\left(r_{p}\right)=2.5\times 10^{-8}$, $r_{p}=2\times
10^{16}$ cm, and $\Delta\log r=1.65$. This is the main function we use to
describe the HC3N abundance in the regions of IRC+10216 probed by our
observations, and an additional narrow component peaking at $r=3\times
10^{16}$ cm with $f=7\times 10^{-7}$ is included to reflect radii larger than
$10^{16}$ cm. This is smaller than the peak derived in Agúndez et al. (2017)
by a factor of 1.4, but this inconsistency is expected since a smaller mass-
loss rate is used in that work. The retrieved abundance profile is shown in
Fig. 7a. along with the results of a chemical model for IRC+10216, which is
described in Section 4.1.

Figure 7: Radial abundance with respect to H2 (top) and volume density
(bottom) profiles of HC3N in the envelope of IRC+10216. Solid black line
represents the best fit profile resulting from the visibility sampled LIME
radiative transfer models, while the dashed gray line is the prediction of a
simple kinetic chemistry model (described in Section 4.1). The shaded red
region represents the radii where the HC3N abundance is best sampled by the
observations.

##  4 Discussion

When compared with the results of our own chemical model as well as the
abundances in Agúndez et al. (2017), the measured HC3N abundances demonstrated
in Fig. 7 is about one order of magnitude larger than predictions of kinetic
chemistry at radii less than $10^{16}$ cm. The disparity with Agúndez et al.
(2017) is actually greater than that since a lower mass-loss rate was adopted
in that work, which will naturally yield higher abundances necessary to
reproduce a given observation.

One factor that could potentially affect abundance calculations in the
envelope of IRC+10216 is the variability of rotational molecular lines with
stellar phase. Pardo et al. (2018) observed periodicity in 3 mm radical,
polyyne, and cyanopolyyne species (including HC3N) that can be either in or
out of phase with the ${\sim}$640 day periodicity of the AGB star as seen in
the infrared. He et al. (2017) observed the same behavior in higher frequency
lines, including the $J=28-27$ blend mapped in this work. We do not expect
this to have a dramatic effect on our results, mainly because the amplitude of
these oscillations are relatively small (ranging from 8%–15% of the average
integrated line flux), but also because the combined data we use for the
$J=28-27$ line was taken, and thereby averaged, over the course of five
months, almost a quarter of the full period. Nevertheless, if we extrapolate
the light curve from He et al. (2017) for the periodicity to the dates of the
ALMA observations (Nov. 2019 – Mar. 2020), we find that the map in Fig. 3a
represents phases near an expected minimum of this line. The $J=38-37$ image
was made from a single observation in May 2018, meaning it would instead be
expected to near the average brightness if it follows the same function. Once
again, however, this is likely a very minor effect on the retrieved abundances
since the amplitude of variation is small, especially when taken in comparison
to the brightness variation from the asymmetric 3D nature of the outflow
(illustrated by the shaded curves in Fig. 6), which is the primary source of
uncertainty in our analysis.

With this in mind, our results indicate that some process is enhancing the
photochemical production of HC3N in the envelope of IRC+10216. Therefore, our
observations of HC3N in the inner layers of IRC+10216 follow a pattern of
previous results for product molecules in this envelope, namely the studies of
CH3CN, NaCN, H2O, and C4H2 (Agúndez et al., 2015; Quintana-Lacaci et al.,
2017; Agúndez et al., 2010; Fonfría et al., 2018). The former is driven by the
radiative association of photoionized CH3+ with HCN, while NaCN is thought to
form via a reaction between NaCl and CN (Agúndez et al., 2008; Petrie, 1996).
H2O and C4H2 have formation routes through reactions with the
photodissociation products C2H and neutral O. HC3N formation in the extended
CSE is likewise governed by the following reactions:

$\mathrm{HCN}+h\nu\rightarrow\mathrm{CN}+\mathrm{H}$ (3)
$\mathrm{C_{2}H_{2}}+\mathrm{CN}\rightarrow\mathrm{HC_{3}N}+\mathrm{H}$ (4)

In common between all these reaction pathways is the need for a source of UV
photons to provide the necessary reactants. The anomalous overabundances for
all four of these species thus suggest a larger flux of UV radiation than
expected in the inner envelope. Agúndez et al. (2015) investigated the
possibility that CH3CN abundances result from the asymmetric clumpiness of the
outflow allowing ionizing photons from the ISRF to penetrate farther into the
outflow than the smooth spherical case; however, this was also not enough to
reproduce their observations. Here, we explore an alternate scenario where the
missing UV is provided internally from a binary companion star orbiting
IRC+10216 at a close distance.

Evidence of such a companion existing nearby IRC+10216 has grown over recent
years, primarily through spatial studies of CO and its isotopologues.
Cernicharo et al. (2015a) and Guelin et al. (1993) note the 16″(${\sim}$700
yr) spacing of the outermost bright gas shells which is neither consistent
with the approximate timescales of thermal pulses (${\sim}10^{4}$ yr) nor the
pulsations exhibited by AGB stars (${\sim}1$ yr). These authors suggest that
the shells could instead be explained by a companion on an elliptical orbit
periodically affecting the mass-loss rate of IRC+10216 during periastron.
Guélin et al. (2018) found that the spacing of these shells decreases moving
inwards, indicating mass transfer or additional dynamics occurring. In
addition, Decin et al. (2015) detected signatures of spiral arm structure at
arcsecond scales which are consistent with a main-sequence 1M☉ companion
orbiting at ${\sim}20$ AU with an orbital period of 55 yr. Lastly, Kim et al.
(2015a) and Kim et al. (2021) detected a secondary point source of optical
light about 0.5″east of IRC+10216 using HST monitoring. Based on its spectral
index, they argue that the object could be a binary companion of spectral type
M1, but this result has yet to be confirmed. In proposed scenarios of Kim et
al. (2015a) and Decin et al. (2015), the binary separation is less than the
minimum beam size of our observations, and in the case of Guélin et al.
(2018), the proposed separation would be ${\sim}1"$. So we expect the region
of the possible companion to be unresolved, or at most very weakly resolved,
from the central AGB star in our observations.

While the exact nature of the orbital system IRC+10216 appears to be a part of
is still not well-characterized due to uncertainties in the above works, a
binary scenario seems very likely. Although IRC+10216 itself is likely far too
red to produce enough UV flux to drive photochemistry in the inner envelope, a
companion with a higher surface temperature could accomplish this. To test
this scheme in comparison to our measurements of HC3N, we utilize a detailed
chemical model of IRC+10216.

###  4.1 Chemical modelling

The chemical kinetics model is that of Van de Sande & Millar (2022), which
includes the effect of stellar companion UV photons on the gas-phase chemistry
throughout the outflow. Several assumptions go into the model to reduce the
complexity of the problem. The one-dimensional model does not include dust
formation and growth, but assumes dust is present from a radius
$R_{\mathrm{dust}}$ onwards. The companion is assumed to be at the center of
the outflow and lie within the dust-free region before $R_{\mathrm{dust}}$.
Considering the scale of the outflow, we consider the effect of misplacing the
companion within the dust-free region is negligible. The photon flux of the
companion is approximated by blackbody radiation and is set by the stellar
radius, $R_{*}$ and blackbody temperature, $T_{\mathrm{comp}}$. The one-
dimensional nature of the model implies that orbital motion cannot be taken
into account. Hence, the companion’s radiation field is assumed to be always
present. This assumption is explored further in Appendix A. The model assumes
the companion to be one of three types: a red dwarf at 4000 K, a solar-like
star at 6000 K or a white dwarf at 10 000 K.

The effects of a clumpy outflow on the chemistry are taken into account using
the porosity formalism, as described by Van de Sande et al. (2018). The
formalism divides the outflow in a stochastic two-component medium, consisting
of a rarified interclump component and an overdense clumped component, and
allows us to include the influence of a clumpy outflow on the penetration of
UV photons along with the effect of local overdensities. Dust and gas are
assumed to be well-mixed in the interclump and clumped components. The
specific locations of the clumps cannot be specified. Rather, the clumpiness
of the outflow is determined by three general parameters: the interclump
density contrast, $f_{\mathrm{ic}}$ which describes the distribution of
material between the clump and interclump component, the volume filling
factor, $f_{\mathrm{vol}}$ which sets fraction of the total outflow volume
occupied by clumps, and the size of the clumps at the stellar surface,
$l_{*}$. As volume filling factor is assumed to be constant throughout the
outflow and the clumps are assumed to be mass conserving, the clumps expand as
they move away from the AGB star.

Figure 8: The retrieved HC3N abundance (black) together with the chemical
modelling results. Gray, solid: smooth outflow without a stellar companion.
Red, dotted: porous outflow (characterised by $f_{\mathrm{ic}}$ = 0.40,
$f_{\mathrm{vol}}$ = 0.50, and $l_{*}=2\times 10^{13}$ cm) without a stellar
companion. Orange, dashed: smooth outflow with a solar-like companion and dust
condensation radius $R_{\mathrm{dust}}$ = 5 R∗. Blue, solid: porous outflow
with a solar-like companion and $R_{\mathrm{dust}}$ = 5 R∗. The shaded red
region represents the radii where the HC3N abundance is best sampled by the
observations.

The values of the three porosity parameters are observationally constrained.
The shells and their clumps were resolved in the VLA survey of Keller (2017).
At a distance of about 24″from the star, the angular sizes of the clumps were
found to range from 2.0″$\times$ 3.0″up to 3.0″$\times$ 7.7″. Assuming a
distance of 130 pc, this corresponds to an upper limit on the physical sizes
of the clumps of $3.9-15.6\times 10^{15}$ cm at a distance of $\sim 4.6\times
10^{16}$ cm from the star. This constrains the value of $l_{*}$ to $2-4\times
10^{13}$ cm. The density contrast between the shell and intershell region was
found to lie between 1.5 and 7.6 at a radial distance of $\sim$ 15″ from the
star. This is consistent with the shell-intershell brightness contrast of
$\geq$ 3 retrieved by Cernicharo et al. (2015b). Assuming this density
contrast can be applied to the clumped and interclump components, the values
of $f_{\mathrm{ic}}$ and $f_{\mathrm{vol}}$ are linked following

$\frac{\rho_{\mathrm{cl}}}{\rho_{\mathrm{ic}}}=\frac{1-(1-f_{\mathrm{fvol}})f_{\mathrm{ic}}}{f_{\mathrm{fvol}}\
f_{\mathrm{ic}}},$ (5)

with $\rho_{\mathrm{cl}}$ and $\rho_{\mathrm{ic}}$ the clumped and interclump
densities, respectively. Lastly, we assume that dust grains are purely
absorbing and therefore neglect the scattering of UV photons. Jones et al.
(2013) showed that for small hydrocarbon dust grains, that provide the
extinction at UV wavelengths, their albedo is essentially zero at these
wavelengths. Such small grains are thought to be present in the inner envelope
(Fonfría et al., 2022). Although the Jones et al. models are strictly for
interstellar dust, their results are likely to hold for IRC+10216; thus
scattering can largely be ignored in the inner regions of the outflow.

The physical parameters adopted for the outflow of IRC+10216, together with
those derived for our best-fit chemical model, are listed in Table 2. The
parent species used are those of Agúndez et al. (2012a), but with the initial
abundance of C2H2 and HCN decreased by a factor two. This was done to ensure
agreement with the results from Fonfría et al. (2008), who observed mid-
infrared vibrational lines of these molecules coming from the innermost
regions of the envelope. Halving the initial C2H2 and HCN lies within the
error of the abundances they retrieve for the region closest to the star, from
the stellar photosphere up to 5.2 R∗, which is where our chemical model
starts.

We find that to reproduce the increase in HC3N in the inner region, both
clumpiness and a companion are necessary. The increase is best reproduced
assuming a solar-like companion and a dust condensation radius of 5 R∗
together with an interclump density contrast $f_{\mathrm{ic}}$ = 0.40. Figure
8 shows the retrieved abundance profile along with chemical model results, for
smooth and porous outflows both with and without a solar-like companion. The
value of $R_{\mathrm{dust}}$ corresponds to the results of Fonfría et al.
(2022), who found gradual dust growth up to 5 R∗. While the chemical model
does not include dust growth, we find that this value of $R_{\mathrm{dust}}$
best reproduced the retrieved abundance profile. The values of
$f_{\mathrm{vol}}$ and $l_{*}$ do not have a large influence on the shape of
the abundance profile and are constrained by the observational results of
Keller (2017).

Although a mass-loss rate of $2.7\times 10^{-5}$ M☉ yr-1 was used for the
radiative transfer model to describe the envelope density within a radius of
15″(Guélin et al., 2018), a lower one is needed for the chemistry model since
the entire extended envelope contributes to the extinction of UV photons, and
lower average mass-loss rates have been obtained when considering radii out to
300″(Decin et al., 2011; Teyssier et al., 2006; Cernicharo et al., 2015a).
Therefore, to better fit the attenuation of the extended molecular shell, the
mass-loss rate was changed from $2.7\times 10^{-5}$ M☉ yr-1, as retrieved by
Guélin et al. (2018) to $1.8\times 10^{-5}$M☉ yr-1.

While the porous model with a solar-like companion shows an increased HC3N
abundance before the molecular shell at $\sim$ 15″, the gentle increase in the
retrieved abundance cannot be reproduced. This is because the model cannot
take the gradual dust growth in the inner wind into account (Fonfría et al.,
2022) and starts only after the dust condensation radius, leading to an abrupt
start of photochemistry in the inner wind and subsequently a sharp increase in
HC3N abundance. Additionally, the porosity formalism provides only a one-
dimensional approximation of the three-dimensional clumpy substructure seen in
IRC+10216’s outflow. Nevertheless, it is once again clear that a classical
chemical model of smooth outflow without a stellar companion does not suffice.

Table 2: Physical parameters of the chemical model Parameters adopted from
Agúndez et al. (2012b)
---
Outflow velocity, $v_{\infty}$ | 14.5 km s-1
Stellar radius, $R_{*}$ | 4 $\times 10^{13}$ cm
Stellar temperature, $T_{*}$ | 2330 K
Exponent $T(r)$, $\epsilon$ | 0.5
Modelling results
Mass-loss rate, $\dot{M}$ |  $1.8\times 10^{-5}$ $\mathrm{M}_{\odot}\ \mathrm{yr}^{-1}$
Interclump density contrast, $f_{\mathrm{ic}}$ | 0.50
Clump volume filling factor, $f_{\mathrm{vol}}$ | 0.25 - 0.80
Companion temperature., $T_{\mathrm{comp}}$, | 6000 K
Companion radius, $R_{\mathrm{comp}}$, |  $8.14\times 10^{10}$ cm
Onset of dust extinction, $R_{\mathrm{dust}}$ | 5 $R_{*}$
Start of the model | $1.025\times R_{\mathrm{dust}}$

##  5 Conclusions

We have presented an investigation of warm HC3N emission in the inner layers
of IRC+10216 utilizing multiple archival ALMA observations. Mapping the
$J=28-27$, $J=30-29$, and $J=38-37$ lines of this molecule reveals the most
compact distribution ($r<10^{16}$ cm) of a cyanopolyyne molecule presently
detected toward the outflow of an evolved star, with emission localized to
radii where CH3CN, NaCN, H2O, and C4H2 were previously observed (Quintana-
Lacaci et al., 2017; Agúndez et al., 2015, 2010; Fonfría et al., 2018). The
emitting region is almost completely distinct from previous works that imaged
lower excitation lines of HC3N, allowing for an in-depth characterization of
this molecule throughout the CSE. Using a 1D model of the physical conditions
surrounding IRC+10216, we used visibility sampled LIME simulations to derive
an average abundance of $10^{-8}$ with respect to H2 for HC3N at the radii
probed by these observations. This is about five to ten times larger than
predicted by simple chemical models of IRC+10216, suggesting that
photochemistry is occurring more rapidly than expected in these warmer layers.

To explain this result, we investigated the possibility that the enhancement
of HC3N surrounding IRC+10216 is caused by an embedded binary companion
emitting UV photons in the inner wind and driving photochemistry. Using known
constraints on the physical environment and the abundances of parent
molecules, we ran a specialized chemical kinetics network to model the radial
HC3N abundance profile for a variety of different companion scenarios. We find
that invoking a solar-like companion and a dust condensation radius of 5 R∗
produces a similar abundance feature in the inner wind to that observed, and
conclude that this is an strong explanation for the anomalous abundances of
product molecules observed at small radii toward IRC+10216.

This result is especially intriguing when considering the newfound prevalence
of binary interactions in AGB outflows (e.g. Decin et al., 2020). Since these
systems appear to be common among nearby evolved stars, it could be that
companions embedded in CSEs have a notable chemical affect on the composition
of material returned to the ISM during the AGB and onward. In order to
understand the full extent of this effect in IRC+10216, further millimeter and
IR studies of photochemistry products (specifically their higher energy
transitions) will be crucial. In addition, rigorous molecular line studies of
C- and O-rich evolved stars that show signs of binary shaping could shed
important light on this topic.

##  6 Acknowledgements

We thank the reviewer for their thorough and insightful comments on this
manuscript. This paper makes use of the following ALMA data:
ADS/JAO.ALMA#2019.1.00507.S, ADS/JAO.ALMA #2011.0.00229.S,
ADS/JAO.ALMA#2016.1.00251.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 cooperation with the
Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and
NAOJ. The National Radio Astronomy Observatory is a facility of the National
Science Foundation operated under cooperative agreement by Associated
Universities, Inc. Support for this work was provided to M. A. S. by the NSF
through the Grote Reber Fellowship Program administered by Associated
Universities, Inc./National Radio Astronomy Observatory. M. A. S. also
acknowledges additional support from the Virginia Space Grant Consortium. M.
V. d. S acknowledges support from the European Union’s Horizon 2020 research
and innovation programme under the Marie Skłodowska-Curie grant agreement No
882991. T. J. M gratefully acknowledges the receipt of a Leverhulme Emeritus
Fellowship and the STFC for support under grant reference ST/P000312/1 and
ST/T000198/1.

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##  Appendix A Effect of orbit geometry

The orbital motion of a stellar companion could induce chemical asymmetries in
the outflow. Because the chemical model is one-dimensional, orbital motion
cannot be taken into account. Moreover, the companion star is assumed to be
located at the center of the AGB star, resulting in a continuous irradiation
of the outflow by the companion star. Van de Sande & Millar (2022) argue that
the limitations of the model are reasonable first-order approximations for a
close-by companion with a short orbital period. This assumption should be
investigated in more detail for the IRC+10216 system because the tentative
orbital solutions presented in Cernicharo et al. (2015a) and Decin et al.
(2015) of respectively 700 and 55 yr cannot necessarily be considered
sufficiently short. For a given point at radius $r$ in the outflow on the
orbital plane of the binary system, assuming a circular orbit, the fraction of
the period the companion spends behind the AGB star is given by

$f\approx\frac{2a\arctan\left[\frac{r+a}{a}\frac{R_{*}}{r}\right]}{2\pi a}$
(A1)

where $a$ is the semi-major axis and $R_{*}$ is the radius of the AGB star,
provided $\bm{r\gg R_{*}}$. Using the empirical values from Cernicharo et al.
(2015a) ($a=42.6\ R_{*}$) and Decin et al. (2015) ($a=8.2\ R_{*}$), we obtain
the results shown in Fig. A1. In both cases, the perceived occultation time of
the companion is less than 10% of the total period at all radii larger than
$R_{\mathrm{dust}}$. In the outer envelope, this value then drops to 4% and 1%
for the long and short period solutions, respectively. From this, we would
only expect time-dependent spatial chemistry differences to occur in the
innermost regions of the envelope, and these scales are unfortunately not
probed by our ALMA observations. Of course, other factors, including the
varied dust extinction between the companion and the point as the companion
orbits, the eccentricity, the presence of density enhancements, clumps, arcs,
etc., do complicate the picture but as a first approximation it is clear that
the companion photons are visible for a very large fraction of the orbit even
in the long period solution from Cernicharo et al. (2015a).

The other important detail here is the inclination of the orbit with respect
to the observer. Since we are only mapping the HC3N distribution at the
systemic velocity, we would only expect chemical asymmetries to arise in
regions opposite where the companion crosses the plane of the sky, which is a
small fraction even for slight inclinations. Though the exact orbit is
uncertain, Decin et al. (2015) proposed that the inclination is larger than
$60^{\circ}$ based on observed ^13CO PV diagrams. An inclination this high
would only cause chemical asymmetries (in the plane of the sky) to be
observable if the envelope was observed during very specific points in the
orbit, which is unlikely given the proposed period of 55 – 700 yr (Decin et
al., 2015; Cernicharo et al., 2015a). Between this and the relatively short
fractional occultation period, if a companion is indeed influencing HC3N
chemistry in the envelope of IRC+10216, we do not expect orbit-induced
asymmetries to be prominent or detectable in the ALMA images presented in this
work.

Figure A1: Occultation time of a companion star orbiting with IRC+10216 semi-
major axis $a$ as seen by a point on the orbital plane as a function of
radius. Magenta and green curves correspond to the semi-major axes proposed in
Cernicharo et al. (2015a) and Decin et al. (2015), respectively.