# Carbon chain diversity in L1544 and IRAS 16293-2422: an astrochemical pathfinder study for the SKAO

Lisa Giani, Eleonora Bianchi, Anthony Remijan, Claudio Codella, Giovanni

Sabatini, Linda Podio, Cecilia Ceccarelli, Marta de Simone, Nadia Balucani,

Paola Caselli, et al.

► To cite this version:

Lisa Giani, Eleonora Bianchi, Anthony Remijan, Claudio Codella, Giovanni Sabatini, et al. Carbon chain diversity in L1544 and IRAS 16293-2422: an astrochemical pathfinder study for the SKAO. Monthly Notices of the Royal Astronomical Society, 2025, 544 (4), pp.4043-4061. 10.1093/mn-ras/staf1941. hal-05398985

HAL Id: hal-05398985

https://hal.science/hal-05398985v1

Submitted on 5 Dec 2025

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# Carbon chain diversity in L1544 and IRAS 16293–2422: an astrochemical pathfinder study for the SKAO

Lisa Giani<sup>1,★</sup> Eleonora Bianchi,<sup>1</sup> Anthony Remijan,<sup>2</sup> Claudio Codella<sup>1,3</sup> Giovanni Sabatini<sup>1</sup>,

Linda Podio,<sup>1</sup> Cecilia Ceccarelli<sup>1,3</sup> Marta De Simone<sup>1,4</sup> Nadia Balucani<sup>1,5</sup> Paola Caselli,<sup>6</sup>

Eric Herbst,<sup>7,8</sup> Francois Lique,<sup>9</sup> Silvia Spezzano,<sup>6</sup> Charlotte Vastel<sup>10</sup> and Brett McGuire<sup>11,12</sup>

Accepted 2025 November 3. Received 2025 October 28; in original form 2025 September 16

## ABSTRACT

Astrochemical observations have revealed a surprisingly high level of chemical complexity, including long carbon chains, in the earliest stages of Sun-like star formation. The origin of these species and whether they undergo further growth, possibly contributing to the molecular complexity of planetary systems, remain open questions. We present recent observations performed using the 100-m Green Bank Telescope of the prestellar core L1544, and the protostellar system IRAS 16293–2422. In L1544, we detected several complex carbon-bearing species, including C<sub>2</sub>S, C<sub>3</sub>S, C<sub>3</sub>N, c-C<sub>3</sub>H, C<sub>4</sub>H, and C<sub>6</sub>H, complementing previously reported emission of cyanopolyynes. In IRAS 16293–2422, we detected c-C<sub>3</sub>H and, for the first time, HC<sub>7</sub>N. Thanks to the high spectral resolution, we refine the rest frequencies of several c-C<sub>3</sub>H and C<sub>6</sub>H transitions. We perform radiative transfer analysis, highlighting a chemical difference between the two sources: IRAS 16293–2422 shows column densities 10 to 100 times lower than L1544. We perform astrochemical modelling, employing an up-to-date chemical network with revised reaction rates. Models reproduce the general trends, with cyanopolyyne and polyynyl radical abundances decreasing as molecular size increases, but underestimate the abundances of cyanopolyynes longer than HC<sub>5</sub>N by up to two orders of magnitude. Current models, which include the dominant neutral–neutral formation routes, cannot account for this discrepancy, suggesting that the chemical network is incomplete. We propose that additional ion–molecule reactions are crucial for the formation of these species. Developing a more comprehensive chemical network for long carbon chains is essential for accurately interpreting present and future observations.

**Key words:** astrochemistry – stars: formation – ISM: abundances – ISM: molecules.

## 1 INTRODUCTION

The formation of a solar-type star and its planetary system begins with a prestellar core, a dense (≥10<sup>5</sup> cm<sup>-3</sup>) and cold (~10 K) condensation of gas typically spanning ~10<sup>3</sup>–10<sup>4</sup> au which eventually collapses under its own gravity (e.g. P. J. Benson & P. C. Myers 1989; E. A. Bergin & M. Tafalla 2007; E. Keto & P. Caselli 2008). As the collapse proceeds, a protostar forms in the central region, heating the gas and generating shocks in the surrounding medium (e.g.

S. W. Stahler & F. Palla 2004). A surprisingly rich chemistry has been observed even in the initial stages of this process, marked by the gas-phase detection of complex molecules relevant for the formation of prebiotic species, such as interstellar complex organic molecules (iCOMs, primarily C-bearing species with at least six atoms; e.g. C. Ceccarelli 2004; E. Herbst & E. F. van Dishoeck 2009; C. Ceccarelli et al. 2023, and references therein). The prestellar phase is particularly rich in deuterated species and iCOMs, which are primarily observed at (sub-)millimeter wavelengths (A. Crapsi et al. 2005; P. Caselli & C. Ceccarelli 2012; L. Bizzocchi et al. 2014; C. Vastel et al. 2014, 2019; I. Jiménez-Serra et al. 2016; S. Spezzano et al. 2017; A. Punanova et al. 2018; S. Scibelli

<footnote n="1"><sup>1</sup>INAF, Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, I-50125 Firenze, Italy</footnote>

<footnote n="2"><sup>2</sup>National Radio Astronomy Observatory, Charlottesville, VA 22903, USA</footnote>

<footnote n="3"><sup>3</sup>Univ. Grenoble Alpes, CNRS, IPAG, F-38000 Grenoble, France</footnote>

<footnote n="4"><sup>4</sup>European Southern Observatory, Karl-Schwarzschild Str 2, D-85748 Garching bei München, Germany</footnote>

<footnote n="5"><sup>5</sup>Dipartimento di Chimica, Biologia e Biotecnologie, Via Elce di Sotto 8, I-06123 Perugia, Italy</footnote>

<footnote n="6"><sup>6</sup>Max-Planck-Institut für extraterrestrische Physik (MPE), Giessenbachstrasse 1, D-85748 Garching, Germany</footnote>

<footnote n="7"><sup>7</sup>Department of Chemistry, University of Virginia, Charlottesville, VA 22904, USA</footnote>

<footnote n="8"><sup>8</sup>Department of Astronomy, University of Virginia, Charlottesville, VA 22904, USA</footnote>

<footnote n="9"><sup>9</sup>Université du Rennes, CNRS, IPR (Institut de Physique de Rennes)UMR 6251, F-35000 Rennes, France</footnote>

<footnote n="10"><sup>10</sup>IRAP, Université de Toulouse, CNRS, UPS, CNES, F-31400 Toulouse, France</footnote>

<footnote n="11"><sup>11</sup>Department of Chemistry, Massachusetts Institute of Technology, Cambridge, MA 02139, USA</footnote>

<footnote n="12"><sup>12</sup>National Radio Astronomy Observatory, Charlottesville, VA 22903, USA</footnote>

<footnote>*E-mail: lisa.giani@inaf.it</footnote>
et al. 2024). Furthermore, recent radio observations of the starless core TMC-1, conducted with the Green Bank Telescope (GBT<fnref n="1" />) and the Yebes 40-m<fnref n="2" />, have also revealed an active chemistry of complex carbon-bearing species, including cyclic hydrocarbons such as c-C₉H₈, benzonitrile (c-C₆H₅CN), cyanopyrene (C₁₆H₉CN), and cyanocoronene (C₂₄H₁₁CN) (e.g. B. A. McGuire et al. 2020; J. Cernicharo et al. 2021, 2023; M. A. Siebert et al. 2022; M. L. Sita et al. 2022; M. Agúndez et al. 2023; R. Fuentetaja et al. 2023; A. Remijan et al. 2023; W. G. D. P. Silva et al. 2023; G. Wenzel et al. 2025a,b, and references therein). This peculiar chemistry is still mostly unexplored due to limitations in current radio facilities, but there are hints that it may not be confined to TMC-1 and could be widespread in starless and prestellar cores (A. M. Burkhardt et al. 2021). On the other hand, a dichotomy in chemical composition has been revealed in protostellar sources. This is characterized by two extreme cases: hot corinos and Warm Carbon Chain Chemistry (WCCC) sources. Hot corinos are regions approximately 100 au in size, where temperatures reach at least 100 K, causing dust mantles to sublimate and, in turn, increasing the abundance of iCOMs (C. Ceccarelli 2004; C. Ceccarelli et al. 2023). Conversely, WCCC sources are poor in iCOMs but enriched in unsaturated small carbon chains (such as c-C₃H₂; N. Sakai et al. 2008; N. Sakai & S. Yamamoto 2013; K. Taniguchi, P. Gorai & J. C. Tan 2024) in regions that can extend up to 1000 au, where temperatures exceed 20 K and CH₄ sublimates. The origin of this chemical diversity remains an open question, and the recent discovery of hybrid sources, such as B335, L483, and IRAS 15398–3359, further complicates the picture (M. Imai et al. 2016; Y. Oya et al. 2017; S. K. Jacobsen et al. 2019; Y. Okoda et al. 2023). While a different duration of the prestellar stage or external ultraviolet (UV) illumination have been proposed as potential explanations for the varying chemistry, the limited number of observations of complex carbon chains prevents a comprehensive understanding of this dichotomy (M. Bouvier et al. 2020, 2022). In this work, we present a spectral survey conducted with the GBT in the Ku (13.5–15.4 GHz) and X (8.–11.6 GHz) bands, targeting one of the prototypical prestellar cores, L1544, along with the molecular envelope surrounding representative protostars embedded within hot corinos, specifically IRAS 16293–2422. The subsequent sections are organized as follows: Section 2 details the two sources, Section 3 describes the observations, Section 4 presents the results, and Section 5 discusses the origin of the observed chemical diversity. Finally, Section 6 provides the conclusions.

## 2 THE SAMPLE

We chose as targets two well-studied sources that are prototypical of their evolutionary stage: L1544, a prestellar core, and IRAS 16293–2422, a multiple protostellar system. For both sources, their physical structure is well-known, and they are recognized for hosting a rich chemistry, making them ideal targets for exploring molecular composition at low frequencies.

### 2.1 The L1544 prestellar core

L1544, located in the Taurus molecular cloud complex at a distance of 170 pc (P. A. B. Galli et al. 2019), is one of the most well-studied prestellar cores. Fig. 1 (upper panel) shows L1544 as traced by the

<figure><img src="image_2.png" /><figcaption>**Figure 1.** Overview of the L1544 (upper panel) and IRAS 16293–2422 (lower panel) regions as traced by the Herschel³ continuum maps at 350 µm. The contours correspond to [5,7,9,12,14]σ, where σ is 12 MJy sr⁻<fnref n="1" /> in L1544, and 1.1 MJy sr⁻<fnref n="1" /> in IRAS 16293–2422, respectively. The FoVs of the GBT observations are shown as circles – solid line for Ku band and dashed line for X-band observations (see Section 3). Diamonds identify the continuum (~1.3 mm) peak position in L1544 (D. Ward-Thompson, F. Motte & P. Andre 1999) and the position of the A/B, and E objects in IRAS 16293–2422 (K. A. Kahle et al. 2023). The size of the Herschel beam (25 arcsec) is shown in the bottom left corner, while the scale bar is shown in the bottom right.</figcaption></figure>

Herschel continuum map at 350 µm. Its dense (~10⁶ cm⁻³) and extremely cold (down to ~7 K) central region (~1000 au) exhibits significant CO depletion (P. Caselli et al. 1999, 2022; A. Crapsi et al. 2007). In contrast, the outer regions host various chemical processes that produce iCOMs and carbon chain species, as evidenced by single-dish and interferometric observations (L. Bizzocchi et al. 2014; C. Vastel et al. 2014, 2016, 2018a, 2019; I. Jiménez-Serra et al. 2016; S. Spezzano et al. 2017; P. Hily-Blant et al. 2018; A. Punanova

<footnote n="1">¹https://greenbankobservatory.org/science/telescopes/gbt/</footnote>

<footnote n="2">²https://rt40m.oan.es/</footnote>
et al. 2018; R. G. Urso et al. 2019). More specifically, methanol peak is about 4000 au north of the core centre while emission from c-C₃H₂ and cyanopolyynes peaks in the southern part of the core (L. Bizzocchi et al. 2014; S. Spezzano et al. 2016; Y. Lin et al. 2022; E. Bianchi et al. 2023). This chemical differentiation has been suggested to arise from uneven illumination, influenced by both the asymmetrical 3D structure of the core and its location at the edge of the molecular cloud (S. S. Jensen et al. 2023).

## 2.2 The IRAS 16293–2422 protostellar system

IRAS 16293–2422 (hereafter IRAS 16293) is an archetypical solar-type star-forming region located in the ρ Ophiuchus cloud complex, at a distance *d* = 141 pc (S. A. Dzib et al. 2018). IRAS 16293 is the first star-forming region where a Class 0 protostar has been identified (P. Andre, D. Ward-Thompson & M. Barsony 1993), and it has been the target of several studies at both mm- and sub-mm wavelengths that have revealed its physical structure (see the Herschel continuum map at 350 µm in Fig. 1 (lower panel). IRAS 16293 is associated with a large envelope extending up to 6000 au (A. Castets et al. 2001; N. Crimier et al. 2010), which surrounds a close binary (protostars A1 and A2) separated by 50 au (M. J. Maureira et al. 2020), and a third protostellar object (labelled B) located 600 au (L. G. Mundy et al. 1992; J. K. Jørgensen et al. 2016) from the A1 + A2 system. Within the same molecular cloud lies the prestellar core IRAS 16293 E, located about 90 arcsec east of the protostellar objects (L. Loinard et al. 2001). The inner regions of IRAS 16293 A and B are characterized by the presence of hot corinos, where the gas phase abundances of many molecules, particularly iCOMs, increase by several orders of magnitude compared to those in the surrounding envelope. IRAS 16293 has been the preferential target of several unbiased spectral line surveys dedicated to unveiling the chemical composition of star-forming regions, such as CHESS (Chemical HeRschel Surveys of star forming regions, C. Ceccarelli et al. 2010), TIMASS (The IRAS 16293 Millimeter and Submillimeter Spectral Survey, E. Caux et al. 2011), and PILS (Protostellar Interferometric Line Survey, J. K. Jørgensen et al. 2016). These studies have led to the discovery of a large chemical complexity in the hot corinos around IRAS 16293 A and B (e.g. S. Cazaux et al. 2003; S. Bottinelli et al. 2004; A. A. Jaber et al. 2014; J. K. Jørgensen et al. 2016, 2018; S. Manigand et al. 2021; H. S. P. Müller et al. 2023, and references therein). On the other hand, the outer envelope is cold (~10–30 K), and characterized by the presence of small cyanopolyynes (up to HC₅N; A. Jaber Al-Edhari et al. 2017), and small carbon chains such as c-C₃H₂, C₂H, C₃H, C₄H (E. Caux et al. 2011), primary revealed thanks to single-dish observations.

## 3 OBSERVATIONS

The observations were conducted on the GBT between June 2019 and February 2020, under project codes AGBT19A.048 and AGBT20A.135 (PI: C. Codella). The target sources, L1544 and IRAS 16293–2422, were observed at coordinates α<sub>12000</sub> = 05h04m16.6 and δ<sub>12000</sub> = +25°10′48′.0, and α<sub>12000</sub> = 16h32m22s.72 and δ<sub>12000</sub> = −24°28′34″.3, respectively. Study of the cyanopolyynes towards L1544 has already been presented in E. Bianchi et al. (2023). L1544 and IRAS 16293 were both observed with the Ku band receiver in combination with the VEGAS (VErsatile GBT Astronomical Spectrometer) spectrometer in multiple set-ups to cover the frequency range (13.5–15.4 GHz). The resulting spectral windows had a bandwidth of 187.5 MHz and comprised 131.072 channels, yielding an extremely high spectral resolution of 1.4 kHz

(equivalent to ~50 m s⁻¹ and 30 m s⁻¹ at 9 and 14 GHz, respectively). The spectral resolution of the IRAS 16293 spectrum was downgraded to 2.98 kHz for c-C₃H, C₆H, and HC₉N, while a resolution of 24 kHz was used for HC₇N to increase the signal to noise (S/N) and identify weak lines. L1544 and IRAS 16293 were observed for 6 and 8 h, respectively. L1544 was successively observed in X band for additional 19 h to cover the frequency range (8.0–11.6 GHz). The telescope’s half-power beamwidth (HPBW) ranged from approximately 54 arcsec in the Ku band to 84 arcsec in the X band. In L1544, the HNCO, methanol, and c-C₃H₂ peaks, as reported in Y. Lin et al. (2022), are included within the largest X beam. Pointing and focus were set using as calibrators the sources 1733–1304 and 0530+1331 for IRAS 16293 and L1544, respectively. The observations were carried out using position switching mode, employing an ON-OFF throw position of 1° west. Data reduction was performed using the gbtidl software.⁴Initially, bad scans were identified and flagged, followed by the performance of ON-OFF position switching subtraction. The calibrated scans from each observing session were then noise-weighted averaged. Subsequently, any RFI (radio frequency interference) and artefacts introduced by Doppler correction were identified and removed. A polynomial fit was applied to remove the baseline. Absolute flux calibration was then performed based on our flux calibrator observations, resulting in a calibration uncertainty of 20 per cent. The typical root mean square (rms) noise level is ~5 mK per channel for both the Ku and X spectra. Line analysis was performed using the GILDAS<fnref n="5" /> CLASS package. Line intensities were converted to main beam temperature (T<sub>MB</sub>) using the beam efficiencies of 0.9864 (10 GHz) and 0.959 (15 GHz), and forward efficiency >0.98, reported in the GBT website¹.

## 4 RESULTS

The spectra observed in Ku band for both L1544 and IRAS 16293, and that observed in X band towards L1544, are shown in Fig. 2 and 3, respectively. The present observations have revealed a significant number of complex carbon chain species, including C₂S, C₃S, C₃N, c-C₃H, C₄H, C₆H, HC₃N, HC₅N, HC₇N, and HC₉N. The lines are considered detected when the line peak intensity S/N is at least 3. Table 1 reports the frequency of each transition (MHz), the energies of the upper level E<sub>up</sub> (K), the line strength Sμ² (D²), the peak intensity T<sub>peak</sub> in T<sub>MB</sub> scale (mK) and the velocity integrated line intensity I<sub>int</sub> (mK km s⁻¹). For L1544, we also include in Table 1 the cyanopolyynes lines previously published by E. Bianchi et al. (2023) from the same data set. The most striking result from the Ku-band comparison (13.9–15.41 GHz), where both targets have been observed, is that L1544 exhibits a richer spectrum than the IRAS 16293 envelope (see Fig. 2). More specifically, in IRAS 16293 the only detected molecular species are c-C₃H and, in addition, ortho- and para-formaldehyde (H₂CO). Conversely, in L1544 we detected in Ku band, formaldehyde, c-C₃H, C₆H (with associated hyperfine structure), HC₇N, and HC₉N. In addition, the observations of L1544 performed in X band led to the detection of C₂S, C₃S, C₃N, C₄H, HC₅N, HC₅N, and further lines of C₆H.

<footnote n="4">⁴https://gbtidl.nrao.edu/</footnote>

<footnote n="5">⁵http://www.iram.fr/IRAMFR/GILDAS</footnote>
<figure><img src="image_3.png" /><figcaption>**Figure 2.** Full continuum-subtracted spectra (in main-beam temperature, $T_{MB}$) of the IRAS 16293-2422 envelope (upper panel) and the L1544 prestellar source (lower panel) observed in Ku band (13.9–15.41 GHz). The brightest lines associated with the detected species (see Table 1 for the list of all the detections) are labelled. The L1544 spectrum has a spectral resolution of 1.4 kHz (~30 m s⁻¹), while the IRAS 16293 spectrum has been here smoothed to 12 kHz (~240 m s⁻¹) to increase the S/N of the emission lines. In the L1544 spectrum, the H₂CO line at ~14.5 GHz is observed in absorption.</figcaption></figure>

## 4.1 L1544

In L1544, we detected a single emission line from C₂S and C₃S, 2 lines from C₃N, 11 lines from C-C₃H, 6 lines from C₄H, and 8 lines from C₆H (see Table 1). In addition, the HC₃N, HC₅N, HC₇N, and HC₉N lines detected in L1544 in the present data set have been analysed in the study by E. Bianchi et al. (2023). The spectra of all the detected transitions for each species are shown in Figs 4, 5, 6, and 7. The exceptionally high spectral resolution of the present data set (~50 m s⁻¹ and 30 m s⁻¹ at 9 and 14 GHz, respectively; see Section 3) enables excellent sampling of the line profiles. All the detected species exhibit a double-peaked line profile, blue- and red-shifted by 0.1–0.2 km s⁻¹ with respect to the L1544 systemic velocity (+7.2 km s⁻¹; e.g. M. Tafalla et al. 1998). The red-shifted peak has a higher intensity than the blue-shifted one. Moreover, the spectral profiles of the current data set closely align with those previously observed in GBT K band towards L1544 for C₄H and C₆H, as reported by H. Gupta et al. (2009). This spectral characteristic profile has also been observed for cyanopolyynes (E. Bianchi et al. 2023), supporting the idea that these molecules trace the same gas component. Given the asymmetric line profiles, we calculated the line-integrated intensities by carefully selecting and integrating the emission over the relevant velocity range. For non-detections, a 3σ upper limit on the integrated intensity ($I_{int}$) was estimated as $3\sqrt{N_{ch}}$d$\sigma$, where $N_{ch}$ is the number of channels included in the integration (corresponding to a line width of 0.4 km s⁻¹; E. Bianchi et al. 2023) and d$\sigma$ is the velocity resolution. We included a calibration uncertainty of 20 per cent in the calculation. To estimate the total column densities, $N^{\text{tot}}$, of the detected species, we used the integrated line intensities reported in Table 1. Given the absence of available collisional coefficients for C-C₃H and C₄H, and the detection of only a single line for C₂S, C₃S, C₃N, and C₆H, we performed a local thermodynamic equilibrium (LTE) analysis for these species. We

assume that the emission fills the beam and we verify this assumption for C₆H, the only species having multiple detected transitions both in X and Ku band. In this case, the column densities derived using different transitions in X and Ku band independently are consistent within the uncertainties. We assume that the same is valid also for the other species. We derived column densities assuming excitation temperatures of both 5 and 12 K, and we report the results as ranges of values. These limits are consistent with the excitation temperatures found by E. Bianchi et al. (2023) for the cyanopolyynes using the large velocity gradient (LVG) method. Consequently, throughout this work we assume that the emission is optically thin and in LTE at the temperatures derived by E. Bianchi et al. (2023). For C-C₃H, C₄H, and C₆H we used the hyperfine fitting method available in GILDAS to fit the hyperfine components to the available astronomical data. We confirm that emission is optically thin ($\tau < 0.1$) for all species, in agreement with what assumed in the present analysis. For C₂S, C₃S, C₃N, and C₆H we calculate the critical density of the detected transition using the Einstein and collisional coefficients from the LAMDA (Leiden Atomic and Molecular Database) data base<fnref n="6" /> (F. L. Schöier et al. 2005, and https://home.strw.leidenuniv.nl/~moldata/), and from K. M. Walker, F. Lique & R. Dawes (2018), E. Sahnoun et al. (2020), and M. Lara-Moreno, T. Stoecklin & P. Halvick (2021). For all species, we estimate $n_{crit} \lesssim 8 \times 10^3$ cm⁻³ at 10 K and $n_{crit} \lesssim 1 \times 10^4$ cm⁻³ at 20 K. In order to verify if the LTE assumption is valid, we consider the physical structure of the two sources. For L1544, the temperature is <12 K and the gas density is $\gtrsim 2 \times 10^4$ cm⁻³ within a radius of 7140 au probed by the GBT observations in X band, following the density profile by E. Keto & P. Caselli (2010). The same holds for IRAS 16293, where the density profile by N. Crimier et al. (2010) predicts a temperature of >15 K and gas densities $\gtrsim 4 \times 10^5$ cm⁻³ within a radius of 3807 au probed by the GBT observations in Ku band. For IRAS 16293 we also checked the critical densities for temperature higher than 20 K, which correspond to the inner regions. For $T > 20$ K, the gas density is larger than $1 \times 10^6$ cm⁻³, while critical densities are always lower than $5 \times 10^4$ cm⁻³. The LTE assumption is thus verified for the transitions of C₂S, C₃S, C₃N, and C₆H. However, for other transitions from C-C₃H and C₄H, collisional coefficients are not available. If the lines are sub-thermally populated, the column densities derived under the LTE assumption should be considered as lower limits.

### 4.1.1 The $C_nX$ family: C₂S, C₃S, C₃N, and C-C₃H

We detected C₂S (N = 2–1, J = 1–0) and C₃S (2–1) (see Fig. 4), towards L1544. When considering the velocity-integrated emission and a temperature range between 5 and 12 K, we obtain $N_{C_2S}^{\text{tot}} = 1.1–2.0 \times 10^{13}$ cm⁻², and $N_{C_3S}^{\text{tot}} = 1.5–2.0 \times 10^{12}$ cm⁻². On the other hand, the C₃N (3/2–1/2) line (assuming again temperature in the 5–12 K range) leads to $N_{C_3N}^{\text{tot}} = 2–4 \times 10^{12}$ cm⁻². In all these cases, the column densities derived from the red-shifted peaks are ~2–3 times higher than those derived from the blue-shifted ones, consistent with the fact that the emission mainly originates from the southern region of L1544, as found for cyanopolyynes by S. Spezzano et al. (2017) and E. Bianchi et al. (2023). The derived beam-averaged column densities are reported in Table 2. These values are consistent within a factor 3 with previous measurements obtained using the IRAM 30-m at 3 mm for C₂S, C₃S (C. Vastel et al. 2018b), and C₃N (C. Vastel et al. 2019).

<footnote n="6">⁶https://home.strw.leidenuniv.nl/moldata/</footnote>
<figure><img src="image_4.png" /><figcaption>**Figure 3.** (a): Full spectrum (in main-beam temperature, $T_{MB}$) of the L1544 prestellar source observed in X band (7.8–11.6 GHz). The brightest lines associated with the detected species (see Table 1 for the list of all the detections) are labelled. The spectral resolution is 2.8 kHz (~60 m s⁻¹). (b, c, d): Zooms in to highlight selected lines of C₄H, C₆H, and C₃N, respectively. The horizontal solid lines correspond to 3σ.</figcaption></figure>

For c-C₃H, we detect 11 hyperfine components corresponding to the N = 1,0–1,1 transition (see Fig. 5), allowing us to derive $N_{C_3H}^{tot}$ between $2 \times 10^{12} \text{ cm}^{-2}$ (5 K), and $2 \times 10^{13} \text{ cm}^{-2}$ (12 K). These values are in agreement with those derived by J. G. Mangum & A. Wootten (1990) at ~15 GHz using the GBT 100-m antenna.

### 4.1.2 The polyynyl radicals family: C₄H and C₆H

We detected C₄H thanks to three hyperfine components of the N = 1–0 J = 3/2–1/2 spectral pattern, as well as the N = 1–0 J = 1/2–1/2 triplet (see Fig. 6). C₆H is detected in both X and Ku bands (see Fig. 7). Specifically, four hyperfine components of the J = 7/2–5/2 transition with upper level energy of 0.80 K are detected in X band (~9 GHz), while four hyperfine components of the J = 11/2–9/2 transition with upper level energy of 2.1 K are detected in Ku band (~15 GHz). The line profiles of both C₄H, and C₆H show the characteristic double-peaked structure observed in L1544 using C₂S, C₃S, C₃N, and c-C₃H (Section 4.1), and cyanopolyynes (E. Bianchi et al. 2023), with the red-shifted peak brighter (by a factor ~2) than the blue-shifted one.

From the LTE analysis of the velocity-integrated emission, we derived $N_{C_4H}^{tot} = 1 \times 10^{13} \text{ cm}^{-2}$ and $3 \times 10^{13} \text{ cm}^{-2}$ for temperatures of 5 and 12 K, respectively. On the other hand, $N_{C_6H}^{tot}$ is 0.58–1.79 $\times 10^{12} \text{ cm}^{-2}$, adopting the 5–12 K range. These values are well consistent with previous GBT observations by H. Gupta et al. (2009). In contrast, our current measurement of the C₄H column density, derived using spectral parameters from the CDMS (Cologne Database for Molecular Spectroscopy) catalogue (see Table 1), is an order of magnitude lower than that reported by Z. Nagy et al. (2019) and H. Gupta et al. (2009), who presumably used the JPL (Jet Propulsion Laboratory) catalogue (H. M. Pickett et al. 1998). This discrepancy will be discussed in detail in Section 5.

## 4.2 IRAS 16293

The observations of IRAS 16293 show the detection of c-C₃H, H₂CO, and HC₇N. The transitions of C₂S, C₃S, C₃N, and C₄H

were not covered by the present observations. On the other hand, we provide upper limits for C₆H, and HC₉N. The c-C₃H line peaks at approximately +4 km s⁻¹, consistent with the systemic velocity of the IRAS 16293 envelope (E. Caux et al. 2011). The c-C₃H and HC₇N transitions are single-peaked, associated with low upper level energies ($E_{up} < 25$ K) and the line profiles are narrow, with a FWHM ≤ 1 km s⁻¹. These findings support that the emission likely originates from the cold IRAS 16293 envelope, and not from the hot corino hosted inside the cloud. The high excitation ($E_{up} = 173.5$ K) p-H₂CO(8,2,6–8,2,7) transition peaks at around +2.6 km s⁻¹, the systemic source velocity of the IRAS 16293 B protostar (E. Caux et al. 2011). Conversely, the low-excitation o-H₂CO(1,1–2,1,2) line ($E_{up} = 23$ K) exhibits an absorption dip at the velocity of the envelope. This suggests absorption of the H₂CO(1,1–2,1,2) emission by the cold gas of the envelope. Fig. 8 (upper panel) presents a comparison between the c-C₃H(3/2-3/2) F = 2–2 spectrum and the p-H₂CO (8,2,6–8,2,7) transition with $E_{up} = 174$ K. The differing peak velocities of these two spectral lines indicate emission originating from distinct physical components. For our subsequent analysis, we did not extend our investigation to H₂CO, owing to its unresolved components and extensive prior studies (e.g. E. F. Dishoeck et al. 1995; C. Ceccarelli et al. 2000; L. Loinard et al. 2000; M. H. D. der Wiel et al. 2019). Instead, our focus was directed towards the emission of carbon chains from the cold envelope of IRAS 16293. The spectra of the detected c-C₃H and HC₇N lines are reported in Figs 9 and 10. The spectral parameter of the observed lines, as well as the line integrated intensities used to calculate $N^{tot}$ for each species, are reported in Table 1. We calculate column densities for a temperature in the range 5–20 K following the results obtained by A. A. Jaber et al. (2014) from the analysis of cyanopolyynes in the envelope of IRAS 16293. The resulting $N^{tot}$ are listed in Table 2. Note that, as for the L1544 case, the reported values are beam-averaged.

### 4.2.1 The C_nX family: c-C₃H

We detected four hyperfine components of the N = 1,0–1,1 transition of c-C₃H (see Fig. 9), with upper level energy ($E_{up}$) of
<table><caption>**Table 1.** Observed molecular lines towards L1544 and IRS 16293. Line peak temperatures and velocity integrated intensities are in $T_{MB}$ scale.</caption><thead><tr><th>Species</th><th>Transition</th><th>v<sup>a</sup><br>(MHz)</th><th>E<sub>up</sub><sup>a</sup><br>(K)</th><th>Sμ<sup>2a</sup><br>(D<sup>2</sup>)</th><th>T<sub>peak</sub><sup>b</sup><br>(mK)</th><th>I<sub>int</sub><sup>b</sup><br>(mK km s<sup>-1</sup>)<br>L1544</th><th>T<sub>peak</sub><sup>b</sup><br>(mK)</th><th>I<sub>int</sub><sup>b</sup><br>(mK km s<sup>-1</sup>)<br>IRAS 16293<sup>c</sup></th></tr></thead><tbody><tr><td>C<sub>2</sub>S</td><td>N = 2–1, J = 1–0</td><td>11119.446400(6E-7)</td><td>0.53</td><td>8.3</td><td>566±7</td><td>154±2</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>S</td><td>J = 2–1</td><td>11561.513200(9E-7)</td><td>0.83</td><td>25.9</td><td>177±7</td><td>45±2</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>N</td><td>N = 1–0 J=3/2–1/2 F = 3/2–1/2</td><td>9884.29320(1.6E-5)</td><td>0.47</td><td>5.6</td><td>16±6</td><td>4±1</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>N</td><td>N = 1–0 J=3/2–1/2 F = 5/2–3/2</td><td>9885.89000(1.0E-5)</td><td>0.47</td><td>16.2</td><td>44±7</td><td>10±1</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>N</td><td>N = 1–0 J=3/2–1/2 F = 3/2–3/2</td><td>9886.09420(2.7E-5)</td><td>0.47</td><td>5.2</td><td>≤ 15</td><td>≤ 3</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>N</td><td>N = 1–0 J=3/2–1/2 F = 1/2–1/2</td><td>9886.99370(3.0E-5)</td><td>0.47</td><td>5.0</td><td>≤ 18</td><td>≤ 3</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>N</td><td>N = 1–0 J=3/2–1/2 F = 1/2–3/2</td><td>9888.79480(3.5E-5)</td><td>0.47</td><td>0.4</td><td>≤ 21</td><td>≤ 4</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>N</td><td>N = 1–0 J=1/2–1/2 F = 1/2–1/2</td><td>9911.18670(4.0E-5)</td><td>0.47</td><td>0.5</td><td>≤ 18</td><td>≤ 3</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>N</td><td>N = 1–0 J=1/2–1/2 F = 1/2–3/2</td><td>9912.98770(3.6E-5)</td><td>0.47</td><td>5.0</td><td>≤ 18</td><td>≤ 3</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>N</td><td>N = 1–0 J=1/2–1/2 F = 3/2–1/2</td><td>9913.64600(3.0E-5)</td><td>0.47</td><td>5.2</td><td>≤ 21</td><td>≤ 4</td><td>...</td><td>...</td></tr><tr><td>C<sub>3</sub>N</td><td>N = 1–0 J=1/2–1/2 F = 3/2–3/2</td><td>9915.44700(1.6E-5)</td><td>0.47</td><td>5.6</td><td>≤ 21</td><td>≤ 4</td><td>...</td><td>...</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=1/2–1/2 F = 1–1</td><td>14686.63000(4.0E-6)</td><td>0.71</td><td>3.8</td><td>71±7</td><td>23±1</td><td>21±4</td><td>13±2</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=1/2–1/2 F = 0–1</td><td>14689.71800(1.2E-5)</td><td>0.71</td><td>1.4</td><td>36±6</td><td>7±1</td><td>≤ 24</td><td>≤ 5</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=1/2–1/2 F = 1–0</td><td>14697.69200(8.0E-6)</td><td>0.71</td><td>2.6</td><td>46±7</td><td>14±2</td><td>≤ 24</td><td>≤ 5</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=1/2–3/2 F = 0–1</td><td>14755.39620(1.2E-5)</td><td>0.71</td><td>1.6</td><td>34±8</td><td>11±2</td><td>≤ 21</td><td>≤ 5</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=1/2–3/2 F = 1–2</td><td>14767.70000(8.0E-6)</td><td>0.71</td><td>2.8</td><td>66±7</td><td>15±2</td><td>15±5</td><td>11±1</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=1/2–3/2 F = 2–1</td><td>14812.01000(8.0E-6)</td><td>0.71</td><td>1.7</td><td>43±7</td><td>11±2</td><td>≤ 24</td><td>≤ 5</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=1/2–3/2 F = 1–1</td><td>14829.57100(6.4E-6)</td><td>0.71</td><td>2.0</td><td>43±8</td><td>12±2</td><td>≤ 24</td><td>≤ 5</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=1/2–3/2 F = 1–0</td><td>14(1.0E-5)</td><td>0.71</td><td>0.7</td><td>≤ 24</td><td>≤ 5</td><td>≤ 24</td><td>≤ 5</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=1/2–3/2 F = 2–1</td><td>14877.67500(4.0E-6)</td><td>0.71</td><td>2.0</td><td>41±6</td><td>12±2</td><td>≤ 24</td><td>≤ 5</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=3/2–3/2 F = 2–2</td><td>14893.05100(4.0E-6)</td><td>0.71</td><td>11.1</td><td>211±5</td><td>71±1</td><td>58±9</td><td>40±2</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=3/2–3/2 F = 1–1</td><td>14895.24300(8.0E-6)</td><td>0.72</td><td>5.3</td><td>104±5</td><td>71±1</td><td>27±3</td><td>14±2</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1.0</sub>–1<sub>1.1</sub> J=3/2–3/2 F = 1–2</td><td>14910.62500(4.0E-6)</td><td>0.72</td><td>0.9</td><td>22±6</td><td>5±1</td><td>≤ 21</td><td>≤ 5</td></tr><tr><td>C<sub>4</sub>H</td><td>N = 1–0 J=3/2–1/2 F = 1–0</td><td>9493.06000(2.0E-6)</td><td>0.46</td><td>2.5</td><td>86±7</td><td>19±2</td><td>...</td><td>...</td></tr><tr><td>C<sub>4</sub>H</td><td>N = 1–0 J=3/2–1/2 F = 2–1</td><td>9497.61500(2.0E-6)</td><td>0.46</td><td>7.4</td><td>247±6</td><td>59±1</td><td>...</td><td>...</td></tr><tr><td>C<sub>4</sub>H</td><td>N = 1–0 J=3/2–1/2 F = 1–1</td><td>9508.00500(2.0E-6)</td><td>0.46</td><td>1.9</td><td>69±7</td><td>16±2</td><td>...</td><td>...</td></tr><tr><td>C<sub>4</sub>H</td><td>N = 1–0 J=1/2–1/2 F = 1–0</td><td>9547.96000(2.0E-6)</td><td>0.46</td><td>1.9</td><td>70±6</td><td>20±1</td><td>...</td><td>...</td></tr><tr><td>C<sub>4</sub>H</td><td>N = 1–0 J=1/2–1/2 F = 0–1</td><td>9551.72000(2.0E-6)</td><td>0.46</td><td>1.5</td><td>68±7</td><td>15±1</td><td>...</td><td>...</td></tr><tr><td>C<sub>4</sub>H</td><td>N = 1–0 J=1/2–1/2 F = 1–1</td><td>9562.90500(2.0E-6)</td><td>0.46</td><td>2.5</td><td>93±7</td><td>21±2</td><td>...</td><td>...</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=3/2 F=4–3 1 = e</td><td>9703.5080(1.0E-2)</td><td>0.80</td><td>111.9</td><td>41±6</td><td>7±2</td><td>...</td><td>...</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=3/2 F=3–2 1 = e</td><td>9703.6000(1.0E-2)</td><td>0.80</td><td>82.9</td><td>28±6</td><td>8±2</td><td>...</td><td>...</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=3/2 F=4–3 1 = f</td><td>9703.8350(1.0E-2)</td><td>0.80</td><td>111.9</td><td>46±6</td><td>10±2</td><td>...</td><td>...</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=3/2 F=3–2 1 = f</td><td>9703.9360(1.0E-2)</td><td>0.80</td><td>82.9</td><td>31±6</td><td>6±2</td><td>...</td><td>...</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=1/2 F=4–3 1 = f</td><td>9758.8270(1.0E-2)</td><td>22.68</td><td>134.3</td><td>≤ 18</td><td>≤ 3</td><td>...</td><td>...</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=1/2 F=3–2 1 = f</td><td>9759.1166(9.2E-3)</td><td>22.68</td><td>99.5</td><td>≤ 18</td><td>≤ 3</td><td>...</td><td>...</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=1/2 F=4–3 1 = e</td><td>9786.7810(4.0E-3)</td><td>22.68</td><td>134.2</td><td>≤ 21</td><td>≤ 4</td><td>...</td><td>...</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=1/2 F=3–2 1 = e</td><td>9787.0050(4.0E-3)</td><td>22.68</td><td>99.5</td><td>≤ 18</td><td>≤ 3</td><td>...</td><td>...</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=3/2 F=6–5 1 = e</td><td>15248.2470(5.0E-3)</td><td>2.12</td><td>192.0</td><td>117±5</td><td>23±1</td><td>≤ 18</td><td>≤ 3</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=3/2 F=5–4 1 = e</td><td>15248.3320(5.0E-3)</td><td>2.12</td><td>159.5</td><td>107±5</td><td>19±1</td><td>≤ 18</td><td>≤ 3</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=3/2 F=6–5 1 = f</td><td>15249.0840(5.0E-3)</td><td>2.12</td><td>192.0</td><td>113±5</td><td>19±1</td><td>≤ 18</td><td>≤ 3</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=3/2 F=5–4 1 = f</td><td>15249.1580(5.0E-3)</td><td>2.12</td><td>159.5</td><td>100±5</td><td>20±1</td><td>≤ 18</td><td>≤ 3</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=1/2 F=6–5 1 = f</td><td>15343.1670(5.0E-3)</td><td>24.02</td><td>205.7</td><td>≤ 18</td><td>≤ 3</td><td>≤ 18</td><td>≤ 3</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=1/2 F=5–4 1 = f</td><td>15343.2370(5.0E-3)</td><td>24.02</td><td>170.9</td><td>≤ 18</td><td>≤ 3</td><td>≤ 18</td><td>≤ 3</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=1/2 F=6–5 1 = e</td><td>15371.5050(5.0E-3)</td><td>24.02</td><td>205.7</td><td>≤ 18</td><td>≤ 3</td><td>≤ 18</td><td>≤ 3</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=1/2 F=5–4 1 = e</td><td>15371.5600(5.0E-3)</td><td>24.02</td><td>170.9</td><td>≤ 18</td><td>≤ 3</td><td>≤ 18</td><td>≤ 3</td></tr><tr><td>HC<sub>7</sub>N</td><td>J = 13–12</td><td>14663.9937(2.0E-4)</td><td>4.93</td><td>906.1</td><td>443±7<sup>d</sup></td><td>151±1<sup>d</sup></td><td>9±2</td><td>9±2</td></tr><tr><td>HC<sub>9</sub>N</td><td>J = 24–23</td><td>13944.8310(1.0E-3)</td><td>8.37</td><td>1947.0</td><td>152±6<sup>d</sup></td><td>24±1<sup>d</sup></td><td>≤ 18</td><td>≤ 2</td></tr><tr><td>HC<sub>9</sub>N</td><td>J = 25–24</td><td>14525.8620(1.0E-3)</td><td>9.06</td><td>2028.1</td><td>157±6<sup>d</sup></td><td>25±1<sup>d</sup></td><td>≤ 15</td><td>≤ 2</td></tr><tr><td>HC<sub>9</sub>N</td><td>J = 26–25</td><td>15106.8910(1.0E-3)</td><td>9.79</td><td>2108.9</td><td>166±6<sup>d</sup></td><td>26±1<sup>d</sup></td><td>≤ 15</td><td>≤ 2</td></tr></tbody></table>

Note. a Frequencies and spectroscopic parameters have been provided by S. Yamamoto et al. (1990), S. Yamamoto et al. (1987), and S. Yamamoto & S. Saito (1994) for C<sub>2</sub>S, C<sub>3</sub>S, and c-C<sub>3</sub>H, respectively, and extracted from the Jet Propulsion Laboratory molecular data base (H. M. Pickett et al. 1998, and https://spec.jpl.nasa.gov/). On the other hand, frequencies and spectroscopic parameters for C<sub>3</sub>N (H. Mikami et al. 1989), C<sub>4</sub>H (C. A. Gottlieb et al. 1983), and C<sub>6</sub>H (M. C. McCarthy et al. 1999; C. A. Gottlieb, M. C. McCarthy & P. Thaddeus 2010), are extracted from the Cologne Database for Molecular Spectroscopy (H. S. P. Müller et al. 2001, 2005; C. P. Endres et al. 2016, and http://www.astro.uni-koeln.de/cdms/). b Errors do not include 20 per cent of calibration. Uncertainties on T<sub>peak</sub> are the local rms of the spectra. Upper limits refer to the 3σ level. Line integrated intensities (I<sub>int</sub>) are computed by carefully selecting and integrating the emission over the relevant velocity range. c, . . . , e means that the frequency of the transition is not in the observed frequency range. d Data presented in E. Bianchi et al. (2023).
<figure><img src="image_5.png" /><figcaption>**Figure 4.** C<sub>2</sub>S, C<sub>3</sub>S, and C<sub>3</sub>N spectra (in T<sub>MB</sub> scale) observed towards L1544. The vertical dashed lines mark the ambient LSR velocity (+7.2 km s<sup>-1</sup>; M. Tafalla et al. 1998).</figcaption></figure>

0.71 K. From the LTE analysis of the velocity-integrated emission of the two available brightest lines, we derived $N_{\text{tot}}^{\text{c-C}_3\text{H}} = 1.0-1.4 \times 10^{12} \text{ cm}^{-2}$, and $3.5-6.5 \times 10^{12} \text{ cm}^{-2}$, assuming a gas temperature between 5 and 20 K, respectively.

The line profile associated with the brightest c-C<sub>3</sub>H line (J = 3/2–3/2 F = 2–2) exhibits a red-shifted wing (up to ~5 km s<sup>-1</sup>). This is due to the fact that the GBT beam in X band (54 arcsec) encompasses the bright and extended molecular outflow emerging from the protostellar system. Fig. 11 shows the position of the GBT beam with respect the CO(3–2) outflow obtained using the Atacama Pathfinder Experiment (APEX; R. Güsten et al. 2006)<fnref n="7" /> telescope by K. A. Kahle et al. (2023). The present GBT observations sample a considerable part of the bright red-shifted lobe. This explains the red-shifted wing in the c-C<sub>3</sub>H line (J = 3/2 to 3/2 F = 2–2) spectrum, which is also consistent with the CO(3–2) spectrum extracted by K. A. Kahle et al. (2023) from the position of the red-shifted outflow peak.

We disentangled the contributions to the c-C<sub>3</sub>H emission from the envelope and the outflow spectral wing by fitting the envelope with a Gaussian profile and associating to the outflow the residual emission. We then derived their respective column densities. For the envelope, the integrated intensity is estimated using a Gaussian fit with full width at half-maximum (FWHM) of 0.5 km s<sup>-1</sup> centred on the systemic velocity of +4 km s<sup>-1</sup> (see Fig. 9). We obtained $N_{\text{tot}} = 1-4 \times 10^{12} \text{ cm}^{-2}$, assuming a gas temperature between 5 and 20 K. The residuals are used to estimate the outflow column density which is in the range $2-4 \times 10^{12} \text{ cm}^{-2}$, adopting a temperature range of 20–40 K (A. Jaber Al-Edhari et al. 2017; K. A. Kahle et al. 2023).

For the outflow component, using the H<sub>2</sub> column density derived from the C<sup>17</sup>O column density of $1.9 \times 10^{15} \text{ cm}^{-2}$ measured by K. A. Kahle et al. (2023), assuming an isotopic ratio of <sup>16</sup>O/<sup>17</sup>O = 2798 (K. D. McKeegan et al. 2011), and a CO/H<sub>2</sub> abundance ratio of $2 \times 10^{-4}$, we obtained an estimate of the c-C<sub>3</sub>H abundance with respect to H<sub>2</sub> of about $10^{-10}$.

### 4.2.2 The polyynyl radicals family: C<sub>6</sub>H

Despite the presence of transitions in the Ku band, no C<sub>6</sub>H lines were detected in the observed frequency range. Assuming a gas temperature between 5 and 20 K, we obtain an upper limit on the column density of $N_{\text{C}_6\text{H}} \le 1 - 5 \times 10^{11} \text{ cm}^{-2}$ from the line at 15248.2470 MHz.

### 4.2.3 The cyanopolyynes family: HC<sub>7</sub>N and HC<sub>9</sub>N

Fig. 10 shows that, once smoothed the spectral resolution to ~0.5 km s<sup>-1</sup>, we report, for the first time, the detection of HC<sub>7</sub>N in IRAS 16293. The J = 13–12 transition was detected at 14.663 GHz with a S/N of 4σ (see Table 1). From a LTE analysis of the line, we derive a column density in the range 6–11 × 10<sup>10</sup> cm<sup>-2</sup>, assuming a gas temperature between 5 and 20 K. We also derived an upper limit on the HC<sub>9</sub>N column density of $N_{\text{HC}_9\text{N}} \le 2.9 - 3.2 \times 10^{10} \text{ cm}^{-2}$ from the line at 15106.8910 MHz.

## 4.3 Suggested corrections for selected c-C<sub>3</sub>H, and C<sub>6</sub>H rest frequencies

In both L1544 and IRAS 16293 spectra, we detected several transitions of c-C<sub>3</sub>H and C<sub>6</sub>H at frequencies significantly offset from those reported in the spectroscopic databases, specifically, JPL for c-C<sub>3</sub>H and CDMS for C<sub>6</sub>H. Thanks to the high spectral resolution of the present observations (1.4 kHz), we were able to accurately determine the discrepancies between the observed frequencies ($v_{\text{obs}}$) with the laboratory values reported in the spectral catalogues ($v_{\text{spec}}$). Table 3 reports the observed frequencies for the c-C<sub>3</sub>H, and C<sub>6</sub>H transitions where the discrepancy is larger than twice the spectral resolution. Note that, when observed in both L1544 and IRAS 16293, the frequency shifts are consistent, supporting the reliability of the corrections. We then propose these corrections in future updates of these data bases. The corrected values for c-C<sub>3</sub>H are in agreement with the recent laboratory measurements reported by C. Xue et al. (2025). For C<sub>6</sub>H, we examined the new catalogue reported by A. Remijan et al. (2023), and the transitions considered show the same frequency shifts as those found in our observations.

# 5 DISCUSSION

## 5.1 Comparison with previous observations

**The L1544 prestellar core:** L1544 has long been recognized as a chemically rich environment, particularly abundant in small carbon chains (C<sub>2</sub>S, C<sub>3</sub>S, C<sub>3</sub>N, and c-C<sub>3</sub>H), and cyanopolyynes (e.g. K. Giers et al. 2022; Y. Lin et al. 2022; E. Bianchi et al. 2023). The present GBT observations further enrich this picture, revealing carbon chains such as C<sub>2</sub>S, C<sub>3</sub>S, C<sub>3</sub>N, and c-C<sub>3</sub>H, along with the polyynyl radicals C<sub>4</sub>H and C<sub>6</sub>H, which are known to be key precursors of cyanopolyynes (J.-C. Loison et al. 2014a; L. Giani et al. 2025). C<sub>2</sub>S was first observed by H. Suzuki et al. (1992), with more recent estimates of the column density derived by C. Vastel et al.

<footnote n="7">⁷https://www.apex-telescope.org</footnote>
<figure><img src="image_6.png" /><figcaption>**Figure 5.** c-C₃H spectra (in T<sub>MB</sub> scale) observed towards L1544. The vertical dashed lines mark the ambient LSR velocity (+7.2 km s<sup>-1</sup>; M. Tafalla et al. 1998). When reported, the frequencies (in MHz) are different from those extracted from the JPL catalogue (H. M. Pickett et al. 1998), and refer to the values needed to centre the spectra to the L1544 systemic velocity.</figcaption></figure>

(2018b) and Z. Nagy et al. (2019) using the IRAM 30-m telescope. Our derived column densities are systematically a factor of 2 higher, but remain in agreement with these studies within the associated uncertainty. This may suggest that the IRAM 30-m observations with telescope HPBWs between 24 and 30 arcsec, centred on the dusty peak of L1544, probably miss the extended emission which is instead recovered by the GBT beam. Indeed, our line profiles are consistent with the C<sub>2</sub>S emission map reported by S. Spezzano et al. (2017), which shows an asymmetric distribution with enhanced emission in the south-east region of L1544. C<sub>3</sub>S and C<sub>3</sub>N detections were previously reported by C. Vastel et al. (2018b, 2019) using the IRAM 30-m antenna, finding column densities compatible (within a factor of 2) with the present ones.

The cyclic isomer c-C₃H was first detected in L1544 by J. G. Mangum & A. Wootten (1990) using the NRAO (National Radio Astronomy Observatory) 43m telescope. In our data set, only the cyclic form (c-C₃H) is detected and not the linear one (l-C₃H). J.-C. Loison et al. (2017) analysed the chemistry of c- and l-C₃H in dark clouds (TMC-1 and B1-b), reporting typical c/l abundance ratios of ~5. Thanks to calculations of rovibrational density of states, they showed that, when formed in highly excited states, the isomerization

between the cyclic and linear isomers (i.e. c-C₃H ⇌ l-C₃H) favour the formation of the cyclic form. In addition, they found that the dissociative recombination (DR) of c,l-C₃H<sup>+</sup> leads mainly to cyclic c-C₃H. Although the formation of c-C₃H is generally favoured, the chemistry of the two isomers is tightly coupled. Given their common precursors and the ease of interconversion, c- and l-C₃H likely coexist and trace similar physical regions, and both should be considered in astrochemical modelling. As a matter of fact, our observed line profiles for c-C₃H are consistent with the l-C₃H map obtained by S. Spezzano et al. (2017) with the IRAM 30-m, revealing a more uniform distribution than C₄H and C₂S. The detection and analysis of cyanopolyynes, up to HC₉N, in L1544 was previously reported by E. Bianchi et al. (2023), who derived abundance ratios of HC₅N:HC<sub>7</sub>N:HC<sub>9</sub>N = 1:0.16:0.25. C₄H and C₆H were previously observed in L1544 by H. Gupta et al. (2009) using GBT in the 18-22 GHz frequency range. Although our column density estimate for C₆H is consistent with their results, our value for C₄H is almost one order of magnitude lower, leading to a significantly higher C₆H/C₄H ratio (4-6 per cent compared to 0.5-0.9 per cent). This discrepancy arises because our analysis adopts the revised dipole moment for C₄H from T. Oyama et al. (2020), which is 2.4 times higher than
<figure><img src="image_7.png" /><figcaption>**Figure 6.** C₄H spectra (in T<sub>MB</sub> scale) observed towards L1544. The vertical dashed lines mark the ambient LSR velocity (+7.2 km s<sup>-1</sup>; M. Tafalla et al. 1998).</figcaption></figure>

<figure><img src="image_8.png" /><figcaption>**Figure 7.** C<sub>6</sub>H spectra (in T<sub>MB</sub> scale) observed towards L1544. The vertical dashed lines mark the ambient LSR velocity (+7.2 km s<sup>-1</sup>; M. Tafalla et al. 1998). When reported, the frequencies (in MHz) are different from those extracted from the CDMS catalogue (H. S. P. Müller et al. 2005), and refer to the values needed to centre the spectra to the L1544 systemic velocity.</figcaption></figure>

the earlier value (2.1 D versus 0.9 D) listed in the JPL and CDMS spectral data bases. This leads to C<sub>4</sub>H abundances that are a factor of 5.44 lower than previous estimates, resulting in proportionally higher C<sub>6</sub>H/C<sub>4</sub>H ratios.

**The IRAS 16293 protostellar system:** IRAS 16293 has been a subject of extensive investigations, with studies focusing on both its cold envelope (e.g. G. A. Blake et al. 1994; E. F. Dishoeck et al. 1995; E. Caux et al. 2011) and the protostellar system (e.g. J. K. Jørgensen
<table><caption>**Table 2.** Carbon chains column densities and abundances measured towards L1544 and IRAS 16293. The column densities are estimated assuming a gas temperature between 5 and 12 K in L1544 (E. Bianchi et al. 2023), and between 5 and 20 K in the envelope of IRAS 16293 (A. Jaber Al-Edhari et al. 2017), respectively. The abundances are derived assuming N(H₂) of 2 × 10²² cm⁻² for L1544 and 5 × 10²² cm⁻² for IRAS 16293 (based on Herschel observations, see the text).</caption><thead><tr><th rowspan="2">Species</th><th colspan="2">N<sub>tot</sub> (×10<sup>12</sup> cm<sup>-2</sup>)</th><th colspan="2">Abundances (×10<sup>-10</sup>)</th><th rowspan="2">Reference</th></tr><tr><th>L1544</th><th>IRAS 16293</th><th>L1544</th><th>IRAS 16293</th></tr></thead><tbody><tr><td>C<sub>2</sub>S</td><td>11–20</td><td>–</td><td>5.5–10.0</td><td>–</td><td>1</td></tr><tr><td>C<sub>3</sub>S</td><td>1.5–2.0</td><td>–</td><td>0.7–1.0</td><td>–</td><td>1</td></tr><tr><td>C<sub>3</sub>N</td><td>2.0–3.6</td><td>–</td><td>1.0–1.8</td><td>–</td><td>1</td></tr><tr><td>c-C<sub>3</sub>H</td><td>2.4–17.8</td><td>1.4–8.9</td><td>1.2–8.9</td><td>0.3–1.8</td><td>1</td></tr><tr><td>C<sub>4</sub>H</td><td>13.7–30.9</td><td>–</td><td>6.8–15.5</td><td>–</td><td>1</td></tr><tr><td>C<sub>6</sub>H</td><td>0.58–1.79</td><td>≤ 0.1–0.5</td><td>0.3–0.9</td><td>≤ 0.02–0.1</td><td>1</td></tr><tr><td>HC<sub>3</sub>N</td><td>> 80</td><td>–</td><td>> 40</td><td>–</td><td>2</td></tr><tr><td>HC<sub>5</sub>N</td><td>30–200</td><td>–</td><td>15–100</td><td>–</td><td>2</td></tr><tr><td>HC<sub>7</sub>N</td><td>4–50</td><td>0.06–0.11</td><td>2–25</td><td>0.012–0.022</td><td>1, 2</td></tr><tr><td>HC<sub>9</sub>N</td><td>0.8–16.0</td><td>≤ 0.029–0.032</td><td>0.4–8.0</td><td>≤ 0.0058–0.064</td><td>1, 2</td></tr></tbody></table>

Note. (1) This work, (2) E. Bianchi et al. 2023

<figure><img src="image_9.png" /><figcaption>**Figure 8.** Upper panel: Comparison between the c-C₃H(3/2-3/2) F = 2-2 (black) and p-H₂CO(8₂,₆-8₂,₇) (E_up = 174 K, magenta) line profiles observed towards IRAS 16293. The vertical scale is in T<sub>MB</sub> units. The vertical dashed black line marks the LSR velocity of the envelope (+4.0 km s⁻¹; E. Caux et al. 2011), while the magenta one is for the LSR velocity of the hot gas around the A and B protostars (+2.6 km s⁻¹; E. Caux et al. 2011). Lower panel: Same as the upper panel using the low-excitation (E_up = 23 K) o-H₂CO(1₁,₁-2₁,₂) emission line (divided by a factor 2.5 for sake of clarity).</figcaption></figure>

et al. 2016, 2018). The narrow linewidths (FWHM ≤ 1 km s⁻¹) and low upper level energies of the molecular lines observed with this study, as well as the analysis of H₂CO lines, confirm that the detected emission predominantly arises from cold envelope gas, with minor contamination from the red-shifted outflow (e.g. in c-C₃H; see Section 4.2 and Fig. 9). Our observations allow for the first detection of HC₇N towards IRAS 16293. We also report, for the first time, upper limits for C₆H and HC₉N. Previous detections of c-C₃H and C₄H were reported by E. Caux et al. (2011). The emission is confirmed to originate from the cold envelope due to its narrow line profiles; however, no column density was derived. Emission from C₄H have been reported also by N. Sakai et al. (2009) and J. E. Lindberg, S. B. Charnley & M. A. Cordiner (2016), using the Mopra 22-m and the Kitt Peak 12-m telescopes, respectively. Based on the

<figure><img src="image_10.png" /><figcaption>**Figure 9.** c-C₃H spectra (in T<sub>MB</sub> scale) observed towards IRAS 16293. The spectral resolution has been downgraded to 2.98 kHz (~60 m s⁻¹) to improve the S/N ratio. The curve on the top panel shows the envelope emission (see the text). The vertical dashed lines mark the ambient LSR velocity (+4.0 km s⁻¹; E. Caux et al. 2011). When reported, the frequencies (in MHz) are different from those extracted from the JPL catalogue (H. M. Pickett et al. 1998), and refer to the values needed to centre the spectra to the IRAS 16293 systemic velocity.</figcaption></figure>

C₄H column density from J. E. Lindberg et al. (2016) and the C₆H column density derived in this work, we derive a C₆H/C₄H ratio of approximately 0.05, consistent with values observed in other sources (see Sections 5.2 and 5.3).

Beside the upper limit on CCS column density, N<sup>tot</sup> ≤ 1 × 10¹⁵ cm⁻², set by E. F. Dishoeck et al. (1995) using the Caltech Submillimeter Observatory (CSO) 1.4-m and James Clerk Maxwell Telescope (JCMT) 15-m antennas, and N<sup>tot</sup> ≤ 4 × 10¹³ cm⁻² reported by M. N. Drozdovskaya et al. (2018) with ALMA in the IRAS 16293
<figure><img src="image_11.png" /><figcaption>**Figure 10.** HC<sub>7</sub>N(13–12) spectra (in T<sub>MB</sub> scale) detected towards IRAS 16293. The spectral resolution has been downgraded to 24 kHz (~0.5 km s<sup>-1</sup>) to increase the S/N ratio. The vertical dashed line marks the ambient LSR velocity of the envelope (+4.0 km s<sup>-1</sup>; E. Caux et al. 2011). The curve represents the Gaussian fit, while the vertical solid segments indicate the brightest hyperfine transitions (F = 12–11, 13–12, and 14–13, in order of increasing rest frequency) expected to contribute to the observed profile.</figcaption></figure>

<figure><img src="image_12.png" /><figcaption>**Figure 11.** The red- and blue-shifted outflowing activity in the IRAS 16293 region as mapped by CO(3–2) using APEX observations adapted from K. A. Kahle et al. (2023). First steps and contours correspond to 3σ (2 K km s<sup>-1</sup>) and 8σ, respectively. The stars indicate the positions of the IRAS 16293 A and B protostars (e.g. J. K. Jørgensen et al. 2016). The circle indicates the HPBW of the present GBT observations in the X band (54 arcsec).</figcaption></figure>

B hot corino, no additional observational data are available for CCS, C<sub>3</sub>S, or C<sub>3</sub>N in the IRAS 16293 envelope for a direct comparison with the present data set.

HC<sub>3</sub>N has previously been detected (i) towards protostars using SMA, VLA, and ALMA interferomets by Y.-J. Kuan et al. (2004), C. J. Chandler et al. (2005), and H. Calcutt et al. (2018), respectively, and (ii) sampling large spatial scales by E. F. Dishoeck et al. (1995) using JCMT 15m (beam size 20 arcsec) and A. Jaber Al-Edhari et al. (2017) using IRAM 30-m (beam size 30 arcsec). E. F. Dishoeck et al. (1995) observed three high-excitation transitions (28–27, 27–26, and 26–25), deriving a rotational temperature of ~115 K and a column density of N<sub>HC<sub>3</sub>N</sub><sup>tot</sup> = 2–5 × 10<sup>12</sup> cm<sup>-2</sup>. In contrast, A.

Jaber Al-Edhari et al. (2017) observed 15 transitions in addition to the three lines reported by E. F. Dishoeck et al. (1995), covering also transitions from (9–8) up to (30–29). Based on these data, they derived an excitation temperature of 20 K and a column density of N<sub>HC<sub>3</sub>N</sub><sup>tot</sup> = 3.2–5.6 × 10<sup>11</sup> cm<sup>-2</sup> for the outer envelope. Since the high-excitation lines detected by E. F. Dishoeck et al. (1995) are not expected to arise from the cold envelope, we adopt the values reported by A. Jaber Al-Edhari et al. (2017) for the outer envelope in the following analysis. In addition, A. Jaber Al-Edhari et al. (2017) also reported the detection of HC<sub>5</sub>N in the envelope of IRAS 16293, deriving a column density of N<sub>HC<sub>5</sub>N</sub><sup>tot</sup> = 3.6–7 × 10<sup>11</sup> cm<sup>-2</sup> at 20 K. The HC<sub>7</sub>N column density derived in this work is consistent with the upper limit (N≤ 5.3 × 10<sup>13</sup> cm<sup>-2</sup>) derived by A. Jaber Al-Edhari et al. (2017). This implies no significant beam dilution differences are observed between GBT and IRAM 30-m data, implying an extended emitting region, larger than both telescope beams. The inferred HC<sub>5</sub>N:HC<sub>7</sub>N:HC<sub>9</sub>N ratios in the IRAS 16293 envelope is 1:0.14: <0.28, closely matching those derived for L1544 (E. Bianchi et al. 2023), suggesting similar cyanopolyyne chemistry across different cold environments.

## 5.2 Comparison between prestellar core and protostellar envelope chemistry

To investigate the chemical differentiation between L1544 and IRAS 16293, we compared the column densities of carbon chains and cyanopolyynes in the two sources (Fig. 12). The comparison reveals several key points:

(i) Cyanopolyynes exhibit similar relative column density ratios in both sources (HC<sub>5</sub>N:HC<sub>7</sub>N:HC<sub>9</sub>N = 1:0.16:0.25 in L1544 and 1:0.14: <0.28 in IRAS 16293) yet their absolute column densities are systematically lower in IRAS 16293, possibly related to a chemical difference between the cold molecular envelope of IRAS 16293 and L1544;

(ii) A similar trend is observed for the polyynyl radicals (e.g. C<sub>4</sub>H, C<sub>6</sub>H), consistent with their role as precursors of cyanopolyynes;

(iii) C-C<sub>3</sub>H is the only species with comparable column densities in both sources, within uncertainties;

(iv) No detections of C<sub>2</sub>S, C<sub>3</sub>S, or C<sub>3</sub>N are available in the literature for the IRAS 16293 envelope, preventing a direct comparison for these species.

In Fig. 13, we compare the column densities of polyynyl radicals C<sub>2</sub>N and cyanopolyynes HC<sub>2n+1</sub>N measured in different star-forming regions. The figure shows that, using the updated C<sub>4</sub>H spectral values, a clear exponential decrease in the column densities of polyynyl radicals is revealed. A similar, though shallower, trend is observed for cyanopolyynes, supporting a chemical connection between the C<sub>2n</sub>H and HC<sub>2n+1</sub>N families. Unsaturated carbon chains are believed to form mainly through neutral-neutral and ion-neutral reactions in the gas phase. Formation on grain surfaces is unlikely due to rapid hydrogenation by accreting H atoms (K. Hiraoka et al. 2000; H. Kobayashi et al. 2017; G. Molpeceres & V. M. Rivilla 2022; G. Fedoseev et al. 2025; M. T. Raaphorst, J. Enrique-Romero & T. Lamberts 2025). L. Giani et al. (2025) revised the neutral-neutral reactions of formation of HC<sub>5</sub>N and found that in cold environments (T~10 K) reactions such as C<sub>6</sub>H + N, C<sub>3</sub>N + C<sub>2</sub>H<sub>2</sub>, and HC<sub>3</sub>N + C<sub>2</sub>H are the dominant ones. The observed exponential decrease in the abundances of polyynyl radicals and cyanopolyynes likely reflects a stepwise chemical growth, in which larger species derive from smaller precursors. For example, HC<sub>5</sub>N may originate from
<table><caption>**Table 3.** List of the transitions observed at frequencies ($v_{obs}$) different from those listed in the spectral catalogues ($v_{spec}$), listed in Table 1. The uncertainties on the observed frequencies are assumed to be twice the spectral resolution of the GBT backend, i.e. 2.8 kHz.</caption><thead><tr><th rowspan="2">Species</th><th rowspan="2">Transition</th><th colspan="2">L1544</th><th colspan="2">IRAS 16293</th></tr><tr><th><i>v</i><sub>obs</sub><sup>a</sup><br>(MHz)</th><th><i>v</i><sub>obs</sub> – <i>v</i><sub>spec</sub><br>(kHz)</th><th><i>v</i><sub>obs</sub><br>(MHz)</th><th><i>v</i><sub>obs</sub> – <i>v</i><sub>spec</sub><br>(kHz)</th></tr></thead><tbody><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1,0</sub>–1<sub>1,1</sub> J=1/2–1/2 F = 1–1</td><td>14686.634</td><td>+4</td><td>–</td><td>–</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1,0</sub>–1<sub>1,1</sub> J=1/2–1/2 F = 0–1</td><td>14689.710</td><td>–8</td><td>–</td><td>–</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1,0</sub>–1<sub>1,1</sub> J=1/2–3/2 F = 0–1</td><td>14755.378</td><td>–18</td><td>–</td><td>–</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1,0</sub>–1<sub>1,1</sub> J=1/2–3/2 F = 1–2</td><td>14767.668</td><td>–32</td><td>14767.670</td><td>–30</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1,0</sub>–1<sub>1,1</sub> J=1/2–3/2 F = 2–1</td><td>14812.014</td><td>+4</td><td>–</td><td>–</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1,0</sub>–1<sub>1,1</sub> J=1/2–3/2 F = 1–1</td><td>14829.580</td><td>+9</td><td>–</td><td>–</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1,0</sub>–1<sub>1,1</sub> J=3/2–3/2 F = 1–1</td><td>14895.247</td><td>+4</td><td>14895.253</td><td>+10</td></tr><tr><td>c-C<sub>3</sub>H</td><td>N = 1<sub>1,0</sub>–1<sub>1,1</sub> J=3/2–3/2 F = 1–2</td><td>14910.615</td><td>–10</td><td>–</td><td>–</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=3/2 F=4–3 I = e</td><td>9703.503</td><td>–5</td><td>–</td><td>–</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 7/2–5/2 Ω=3/2 F=4–3 I = f</td><td>9703.840</td><td>+5</td><td>–</td><td>–</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=3/2 F=6–5 I = e</td><td>15248.243</td><td>–4</td><td>–</td><td>–</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=3/2 F=5–4 I = e</td><td>15248.321</td><td>–11</td><td>–</td><td>–</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=3/2 F=6–5 I = f</td><td>15249.087</td><td>+3</td><td>–</td><td>–</td></tr><tr><td>C<sub>6</sub>H</td><td>J = 11/2–9/2 Ω=3/2 F=5–4 I = f</td><td>15249.161</td><td>+3</td><td>–</td><td>–</td></tr></tbody></table>

Note. $^a$ The observed frequencies were derived from the spectra at the native spectral resolution of 1.4 kHz.

<figure><img src="image_13.png" /><figcaption>**Figure 12.** Column densities of carbon chains and cyanopolyynes (see Table 2) observed in L1544 and IRAS 1629. The upper limits on C<sub>6</sub>H and HC<sub>9</sub>N column densities in L1544 are shown with arrows. The arrow for the lower limit on the HC<sub>3</sub>N column density is due to opacity effects (E. Bianchi et al. 2023). The C<sub>4</sub>H and HC<sub>5</sub>N values in IRAS 16293 are taken from J. E. Lindberg et al. (2016) and A. Jaber Al-Edhari et al. (2017), respectively.</figcaption></figure>

<figure><img src="image_14.png" /><figcaption>**Figure 13.** Comparison between column densities of polyynyl radicals C<sub>2n</sub>H (left panel) and cyanopolyynes HC<sub>2n+1</sub>N (right panel) measured in different star-forming regions. Values for L1544 are from this work, and from E. Bianchi et al. (2023) and K. Giers et al. (2023), while values for IRAS 16293 are from this work, and from A. A. Jaber et al. (2014) and J. E. Lindberg et al. (2016). Values for sources different than L1544, and IRAS 16293 are taken from T. Oyama et al. (2020, and references therein).</figcaption></figure>

HC<sub>3</sub>N and C<sub>6</sub>H, which themselves derive from smaller species such as C<sub>4</sub>H.

In general, the formation of carbon chains requires large amounts of atomic carbon. So far, two main scenarios have been proposed: either the gas is chemically young and carbon is not yet locked into CO, or a large irradiation by UV photons or cosmic rays drives CO dissociation, releasing carbon atoms into the gas phase (e.g. M. Agúndez & V. Wakelam 2013; N. Sakai & S. Yamamoto 2013; S. Spezzano et al. 2016, 2017). Large values of $\zeta_{CR}(>10^{-16} s^{-1})$ are however excluded by observations in L1544, as models better reproduces ions abundances when adopting $\zeta_{CR} = 2-3 \times 10^{-17} s^{-1}$ (E. Redaelli et al. 2021). To further explore the origin of the chemical differences between the two sources, we performed astrochemical modelling with varied physical parameters. The model description and results are presented in the following section.

## 5.3 Astrochemical modelling

### 5.3.1 Model description

We ran the MyNahoon astrochemical model (V. Wakelam et al. 2005, 2010) to investigate the chemical diversity between L1544 and the cold molecular envelope of IRAS 16293. The model has been extensively described in previous works (L. Podio et al. 2014; C. Codella et al. 2017; L. Giani et al. 2023). In summary, it calculates the gas phase abundances of ~500 molecules as a function of time for some specific physical parameters, such as gas temperature, T, volume density, $n_H$, cosmic-ray ionization rate, $\zeta_{CR}$, and visual extinction, $A_V$. We adopted the GRETOBAPE gas-phase reaction network (L. Tinacci et al. 2023), including recent updates on the chemistry of formation of HC<sub>5</sub>N (L. Giani et al. 2025) and c-C<sub>3</sub>H (J.-C. Loison et al. 2017). Surface chemistry (including adsorption and
<table><caption>**Table 4.** Physical parameters and initial elemental abundances used in the cold prestellar core/envelope model. The upper half table lists the adopted parameters: H density, $n_H$, temperature, T, cosmic-ray ionization rate, $\zeta_{CR}$, visual extinction, $A_v$, dust grain radius, $a_d$, and grain density, $\rho_d$. The lower half table lists the initial elemental abundances relative to H nuclei, adapted from E. B. Jenkins (2009).</caption><thead><tr><th rowspan="2">Parameter</th><th colspan="2">Physical parameters of the cold gas</th></tr><tr><th>Value</th><th>Parameter</th><th>Value</th></tr></thead><tbody><tr><td>$n_H$ (cm<sup>-3</sup>)</td><td>2 &times; 10<sup>4</sup></td><td>$A_v$ (mag)</td><td>1–10</td></tr><tr><td>$T$ (K)</td><td>10</td><td>$a_d$ (μm)</td><td>0.1</td></tr><tr><td>$\zeta_{CR}$ (s<sup>-1</sup>)</td><td>0.1 – 10 &times; 10<sup>-16</sup></td><td>$\rho_d$ (g cm<sup>-3</sup>)</td><td>3</td></tr><tr><td colspan="4">Initial elemental abundances (wrt H)</td></tr><tr><td>Element</td><td>Abundance</td><td>Element</td><td>Abundance</td></tr><tr><td>$H_2$</td><td>5.0 &times; 10<sup>-1</sup></td><td>$Si^+$</td><td>8.0 &times; 10<sup>-9</sup></td></tr><tr><td>$He$</td><td>9.0 &times; 10<sup>-2</sup></td><td>$P^+$</td><td>2.0 &times; 10<sup>-10</sup></td></tr><tr><td>$C^+$</td><td>2.0 &times; 10<sup>-5</sup></td><td>$Na^+$</td><td>2.0 &times; 10<sup>-9</sup></td></tr><tr><td>$O$</td><td>2.6 &times; 10<sup>-5</sup></td><td>$Fe^+$</td><td>3.0 &times; 10<sup>-9</sup></td></tr><tr><td>$N$</td><td>6.2 &times; 10<sup>-6</sup></td><td>$Cl^+$</td><td>1.0 &times; 10<sup>-9</sup></td></tr><tr><td>$S^+$</td><td>8.0 &times; 10<sup>-8</sup></td><td>$F^+$</td><td>1.0 &times; 10<sup>-9</sup></td></tr></tbody></table>

desorption) is not treated in the network, except for the formation of $H_2$ in dust grains. As unsaturated carbon chains are formed primarily in the gas phase, we do not expect a significant impact from the missing grain-surface processes. The initial elemental abundances and physical parameters adopted in the simulations are summarized in Table 4. To simulate the evolution of a typical translucent cloud into a molecular cloud, we assumed as initial conditions that elements with ionization potentials below 13.6 eV (e.g. C, S, Si, P) are ionized, while those with higher values (N, O) are neutral. In all simulations, hydrogen is assumed to be fully molecular. The 0D model cannot capture the asymmetric spatial distributions of different species, likely arising from the source's physical structure (S. S. Jensen et al. 2023). However, we account for this limitation by comparing beam-averaged abundances with those predicted by the model. The key novelty of this model lies in the updated chemical network with thoroughly studied reaction rates, making it a valuable tool for exploring how the abundances of long carbon chains vary with different parameters. To assess the influence of cosmic rays and UV radiation on carbon chain formation, we ran a grid of models assuming the following ranges: (i) $\zeta_{CR}$ between $1 \times 10^{-17}$ s$^{-1}$ and $1 \times 10^{-15}$ s$^{-1}$ (C. Ceccarelli et al. 2014; E. Redaelli et al. 2021, 2025; G. Sabatini, S. Bovino & E. Redaelli 2023; A. Socci et al. 2024); (ii) $A_v$ between 1 and 10 mag; and (iii) C/O elemental abundance ratios between 0.77 and 1.1. Since elemental depletion in the outer envelope of IRAS 16293 may not yet be complete, we also ran a model in which the initial elemental abundances from Table 4 were increased by a factor of 10. In order to compare the observed column densities with the abundances predicted by the model, we consider the $H_2$ column density to be at least $2 \times 10^{22}$ cm$^{-2}$ and $5 \times 10^{22}$ cm$^{-2}$ for L1544 and IRAS 16293, respectively, from Herschel observations (J. M. Kirk et al. 2013; S. Spezzano et al. 2016; B. Ladjelate et al. 2020) and in agreement with previous estimates (E. F. Dishoeck et al. 1995; A. Crapsi et al. 2005). We note that these values correspond to the integrated $H_2$ column density along the line of sight, while the emission of carbon chains and cyanopolyynes may originate from an external layer, where the $H_2$ column density is likely lower. No direct measurements of $H_2$ in these region are available, making it difficult to accurately determine molecular abundances. To avoid assuming an arbitrarily reduced $H_2$ column density, we adopted the integrated values while emphasizing that the resulting abundances reported in

Table 2 should be used with caution. Accordingly, in the discussion of the chemical modelling, we focus primarily on abundance ratios rather than absolute abundances.

### 5.3.2 Abundance of carbon chains

As a first test, we examined how variations in $A_v$, $\zeta_{CR}$, and C/O elemental ratio affect the predicted abundances of carbon chains. For these models, we adopted the same gas density and temperature for L1544 and IRAS16293, given their similar physical conditions. Fig. 14 compares the observed and modeled abundances of $HC_5N$, $HC_7N$, $C_6H$, and $c-C_3H$ under these varying conditions. For both sources the best agreement with observations is obtained for cosmic-ray ionization rate $\zeta_{CR} = 1 \times 10^{-17}$ s$^{-1}$, visual extinction $A_v = 10$ mag and C/O = 0.77. In all models, cyanopolyynes are under-produced by approximately 2 orders of magnitude compared to L1544 values. Conversely, good agreement is achieved for $C_6H$ and $c-C_3H$ abundances at times around $2–3 \times 10^5$ yr. The recent revision of the gas-phase formation reactions of $HC_5N$ by L. Giani et al. (2025) suggested that the $C_6H + N$ reaction is the most important in cold regions. However, while the model considering this updated network correctly reproduces the $C_6H$ abundance, it fails to reproduce that of $HC_5N$. If $C_6H$ were indeed the major $HC_5N$ precursor, the model should accurately predict $HC_5N$ abundances. This discrepancy strongly suggests that additional formation reactions are missing from the chemical network, not only for $HC_5N$ but also for the longer cyanopolyynes. In IRAS 16293, the observed abundances of $c-C_3H$ and $HC_5N$ are well reproduced by the model at an age of approximately $7 \times 10^5$ yr, in line with the source being more chemically evolved than L1544. At this time, however, the model underestimates the abundance of $HC_7N$ by more than 1 order of magnitude. Unfortunately, the upper limit on the $C_6H$ abundance is not sufficiently constraining for a meaningful comparison. When adopting less depleted initial elemental abundances, all species show comparable abundances at $10^5$ yr (with variations of less than one order of magnitude). At later times ($t > 10^6$ yr), however, their abundances decrease by several orders of magnitude relative to the predictions obtained with the more depleted abundances, thereby worsening the agreement with the observations.

### 5.3.3 Abundance ratios of carbon chains

To overcome the substantial uncertainties associated with observed $H_2$ column densities, we also compare the model predictions with observed abundance ratios of the different species. Fig. 15 shows the comparison of the predicted $C_6H/C_4H$ ratios with those derived from observations, using both the old and revised dipole moment of $C_4H$. As discussed in Section 5.2, T. Oyama et al. (2020) estimated the $C_4H$ dipole moment to be 2.4 times larger than the previous value, implying that previous estimates of its column density were overestimated by a factor of ~6. When the corrected dipole moment is used, the observed $C_6H/C_4H$ ratios, including those reported in this work, are in agreement with model predictions within a factor 3. This possible discrepancy highlights the need for a revision of the formation and destruction pathways for these species.

Fig. 16 compares the predicted $HC_5N/HC_3N$, $HC_7N/HC_5N$, and $HC_9N/HC_7N$ ratios with those derived from the observations presented in this work. There is a clear discrepancy between models and observations with the exception of the $HC_5N/HC_3N$ ratio, which is reasonably reproduced at $t \sim 2–3 \times 10^5$ yr with $A_v = 10$ mag.
<figure><img src="image_15.png" /><figcaption>**Figure 14.** Comparison between the observed and modeled abundances of HC₅N, HC₇N, C₆H, and c-C₃H for different ζCR, A_v, and C/O ratios. Panels (a, b, c) show the predicted abundances of HC₅N and HC₇N as a function of the parameters indicated in the legends. Panels (d, e, f) and (g, h, i) show the same for C₆H and c-C₃H, respectively. Observed values from L1544 and IRAS 16293 are shown as shaded bands, with L1544 indicated by dashed hatching and IRAS 16293 by a solid (non-hatched) shading. We assumed an N(H₂) of 2×10²² cm⁻² for L1544 and 5×10²² cm⁻² for IRAS 16293 (based on Herschel observations, see the text).</figcaption></figure>

Fig. 17 further illustrates that this discrepancy is not limited to a single source: even across different regions with different evolutionary ages, current models consistently underestimate the HC₇N/HC₅N and HC₉N/HC₇N ratios. Specifically, the predicted HC₇N/HC₅N ratios are lower than observed by factors of ~2–10, while the HC₉N/HC₇N ratios are underestimated by factors of ~7–50. Although existing models do reproduce the general trend of decreasing abundance with increasing molecular size, they fail to reproduce the observed abundance ratios, particularly for the longest chains. In order to explore the possibility that cyanopolyynes and carbon chains arises from a UV illuminated region, as proposed by S. Spezzano et al. (2016) for L1544, we verified the impact of different A_v values on the model predictions. Fig. 18 shows the predicted abundances (with respect to H₂) of various carbon chains as a function of the visual extinction A_v. We conservatively assume a gas volume density nH₂ = 10⁴ cm⁻³. Using the current network, cyanopolyynes and polyynyl radicals show a strong decrease by 1 to 4 orders of magnitude from high to low A_v. In contrast, the abundance of c-C₃H remains nearly constant across the entire A_v range (between 1 and 10). Given that the major formation routes for cyanopolyynes in the current chemical network are neutral-neutral reactions, it is not surprising that our model predicts a decrease in cyanopolyyne abundance instead of an increase for higher UV illumination. On

the contrary, the model successfully explains the presence of small carbon chains, such as c-C₃H and c-C₃H₂, in highly irradiated environments (C. Favre et al. 2018). This suggests that additional formation reactions for cyanopolyynes and their precursors are missing in the network. More specifically, additional ion-molecule reactions could explain the enhancement of cyanopolyynes in UV exposed environments.

### 5.3.4 Observations versus models: carbon chains formation routes

Following the comparison of model predictions and observations reported in the previous sections, our main findings are:

(i) Models reproduce c-C₃H and C₆H in L1544 for A_v = 10, standard ζCR and C/O values, and times of ~2–3 × 10⁵ yr. For the same conditions, cyanopolyynes are underproduced by a factor ~100. For IRAS 16293, c-C₃H, and HC₅N are reproduced by the model for A_v = 10 and times of ~7 × 10⁵ yr, but not HC₇N.

(ii) Models reasonably reproduce the abundance ratios between species within the same chemical family for A_v = 10 and evolutionary time-scales in the range 2–7 × 10⁵ yr. The C₆H/C₄H and HC₅N/HC₃N ratios are reproduced in both L1544 and IRAS 16293 (see Figs 15 and 16), while the HC₇N/HC₅N and HC₉N/HC₇N ratios
<figure><img src="image_16.png" /><figcaption>**Figure 15.** Comparison of the C<sub>6</sub>H/C<sub>4</sub>H ratio (R<sub>N</sub>) obtained for L1544 with those derived in cold molecular clouds and protostars. Open circles correspond to data from H. Gupta et al. (2009, and references therein), while filled circles correspond to the same data corrected taking into account the revised dipole moment of C<sub>4</sub>H (see the text in Section 5 and T. Oyama et al. 2020). The R<sub>N</sub> value derived in this work for L1544 and IRAS 16293 are also shown. The temperatures at which the column densities of C<sub>4</sub>H and C<sub>6</sub>H have been derived are reported next to each point. R<sub>N</sub> values predicted by the model at times of 2×10<sup>5</sup> and 7×10<sup>5</sup> yr are shown as solid and dashed lines, respectively. Model predictions are obtained assuming ζ<sub>CR</sub> = 10<sup>-17</sup> s<sup>-1</sup> and A<sub>v</sub> = 10 mag.</figcaption></figure>

<figure><img src="image_17.png" /><figcaption>**Figure 16.** Comparison between predicted and observed abundance ratios of HC<sub>5</sub>N/HC<sub>3</sub>N, HC<sub>7</sub>N/HC<sub>5</sub>N, and HC<sub>9</sub>N/HC<sub>7</sub>N for different visual extinctions (A<sub>v</sub>). Model predictions for HC<sub>5</sub>N/HC<sub>3</sub>N, HC<sub>7</sub>N/HC<sub>5</sub>N and HC<sub>9</sub>N/HC<sub>7</sub>N are shown as lines: solid lines correspond to A<sub>v</sub> = 10, while dashed lines to A<sub>v</sub> = 3. The observations, based on values derived for both L1544 and IRAS 16293 in this work, are represented as shaded bands. Model predictions are obtained assuming ζ<sub>CR</sub> = 10<sup>-17</sup> s<sup>-1</sup> and A<sub>v</sub> = 10 mag.</figcaption></figure>

are underestimated by a factor of ~2–50 likely due to the incomplete chemical network for large cyanopolyynes (see Fig. 17);

(iii) Cyanopolyynes are predicted to be abundant in young objects, while their abundances decrease at later evolutionary stages (t > 3 ×10<sup>5</sup> yr; see Fig. 14). As proposed for the starless core TMC-1, this behaviour can be explained by the larger availability of carbon atoms before being converted into CO (M. Agúndez & V. Wakelam 2013). This explains the higher cyanopolyyne abundances observed in L1544 compared to IRAS 16293, since L1544 is a younger source.

<figure><img src="image_18.png" /><figcaption>**Figure 17.** Comparison of the HC<sub>7</sub>N/HC<sub>5</sub>N (left panel), and HC<sub>9</sub>N/HC<sub>7</sub>N (right panel) ratios obtained for L1544 with those derived in cold molecular cloud. Values for L1544 are from this work, while values for sources different than L1544 are taken from E. Bianchi et al. (2023, and references therein). Values predicted by the model at times of 7×10<sup>4</sup>, 2×10<sup>5</sup> and 7×10<sup>5</sup> yr are shown as solid, dashed and dotted lines, respectively. Model predictions are obtained assuming ζ<sub>CR</sub> = 10<sup>-17</sup> s<sup>-1</sup> and A<sub>v</sub> = 10 mag. Sources for which both the HC<sub>7</sub>N/HC<sub>5</sub>N and HC<sub>9</sub>N/HC<sub>7</sub>N ratios are available are in boldface. Labels indicate the regions where sources are located: Taurus, Lupus, or other regions (Chameleon, Cepheus, and Aquila Rift).</figcaption></figure>

<figure><img src="image_19.png" /><figcaption>**Figure 18.** Predicted abundances (with respect to H<sub>2</sub>) of c-C<sub>3</sub>H (red solid line), C<sub>4</sub>H (purple dotted line), C<sub>6</sub>H (cyan dotted line), HC<sub>3</sub>N (grey dashed line), HC<sub>5</sub>N (green dashed line), HC<sub>7</sub>N (blue dashed line), and HC<sub>9</sub>N (orange dashed line) as a function of A<sub>v</sub>. The model assumes n<sub>H2</sub> = 10<sup>4</sup> cm<sup>-3</sup>, T = 10 K, and ζ<sub>CR</sub> = 10<sup>-17</sup> s<sup>-1</sup>. The predictions refer to the model at 10<sup>5</sup> yr.</figcaption></figure>

(iv) Current models do not support the hypothesis that cyanopolyyne abundances are enhanced by elevated cosmic-ray ionization rate or intense UV irradiation, despite indications from the asymmetric spatial distribution of carbon chains and cyanopolyynes observed towards the south-eastern region of L1544 (S. Spezzano et al. 2016; E. Bianchi et al. 2023). The discrepancy between predicted and observed abundances is particularly pronounced for cyanopolyynes. This highlights the incompleteness of current chemical networks, especially regarding their formation reactions, which in turn impedes a comprehensive understanding of their primary formation mechanisms.

(v) Although previous studies have shown that neutral–neutral reactions are the primary formation routes of cyanopolyynes (L.
Giani et al. 2025), they cannot fully explain the observed abundances, especially in L1544. We therefore strongly suggest that additional ion-molecule reactions, beyond those currently included in the networks, are involved, as the production of ionic species can be enhanced in highly irradiated environments. A thorough revision of the formation pathways of the $HC_5NH^+$ ion, which could be a major precursor of $HC_5N$, is needed to verify this hypothesis (see Appendix for more details).

## 5.4 Future impact of SKA observations on carbon chain chemistry

The present observations confirm that complex carbon chemistry is already active from the earliest stages of Sun-like planetary system formation, with these species potentially contributing to the synthesis of molecules of prebiotic relevance. However, single-dish facilities have typically an angular resolution of 54–84 arcsec in the 8–15 GHz range (corresponding to ~10 000 au scales for nearby star-forming regions), which allow for the observation of these species in extended sources, such as prestellar cores or protostellar envelopes, leaving it unclear whether they also exist at smaller angular scales. New possibilities to observe these species at high-angular resolution will be offered by next-generation facilities such as MeerKAT<fnref n="8" /> and the Square Kilometre Array Observatory (SKAO)<fnref n="9" />. Specifically, MeerKAT Band 5b (covering 8.3–15.4 GHz), currently under development, will enable observations of prestellar cores, protostellar envelopes, and protostellar shocks down to an angular resolution of a few hundreds au in the nearest star-forming regions. Subsequently, SKA-Mid Band 5 (covering the 4.6–15.4 GHz range) will allow for the exploration of this chemical complexity down to the 100 au scale within planet-forming regions. With these new facilities, multiple transitions of cyanopolyynes up to $HC_{11}N$, and their precursors $C_4H$ and $C_6H$, will be detectable in several sources, resolving emission both spatially and spectrally in just a few hours of observation for extended sources such as prestellar cores and protostellar shocks, and in 100–1000 h for compact planet-forming discs. These future SKA observations, coupled with an exhaustive revision of the chemical reaction networks, will provide crucial new insights into the formation and destruction routes of these species and their possible role in promoting prebiotic chemistry.

## 6 CONCLUSIONS

In this paper, we report new observations of several carbon chains conducted with the 100-m GBT. These observations include Ku-band data (13.5–15.4 GHz) targeting both the L1544 prestellar core and the IRAS 16293 protostellar envelope, along with additional X-band data (8.0–11.6 GHz) towards L1544. The main results of the paper are:

(i) In L1544, we detect emission lines from $C_2S$, $C_3S$, $C_3N$, c-$C_3H$, $C_4H$, and $C_6H$ (with $E_{up} \le 2$ K), in addition to those of $HC_3N$, $HC_5N$, $HC_7N$, and $HC_9N$, already reported in E. Bianchi et al. (2023). All the lines have FWHM around 0.4 km s⁻¹ and show similar double-peaked profiles, suggesting a common underlying gas component.

(ii) In IRAS 16293, we detect c-$C_3H$ and, for the first time, $HC_7N$ (emission lines with $E_{up} \le 5$ K). We also report, for the first time, upper limits for $C_6H$ and $HC_9N$. The detected lines have FWHM $\le 1$ km s⁻¹ and they arise from the cold envelope. For c-$C_3H$ emission,

we have contribution from both the cold envelope and the red-shifted outflow that we spectrally disentangle.

(iii) We refine the spectroscopic frequencies of several c-$C_3H$ and $C_6H$ transitions, thanks to the high spectral resolution of the observations (1.4 kHz). The shifts between calculated and observed frequencies are consistent in both L1544 and IRAS 16293.

(iv) We performed LTE analysis to derive column densities of the detected species, assuming gas temperatures in the range 5–12 K for L1544, and 5–20 K for IRAS 16293, respectively. All column density estimates obtained in this work are consistent with previous measurements, where available. The only exception is $C_4H$, for which we used the revised dipole moment reported by T. Oyama et al. (2020), that yields column densities lower by a factor ~6 compared to previous estimates. In IRAS 16293 we also calculated, for the first time, the c-$C_3H$ column density and abundance (~$10^{-10}$) in the outflow wing.

(v) IRAS 16293 is less rich in carbon chains compared to L1544, with all detected species (except for c-$C_3H$) showing column densities lower by a factor ~10–100. Despite this difference, the ratios between polyynyl radicals ($C_{2n}H$) and cyanopolyynes ($HC_{2n+1}N$) are similar in both sources (i.e. $C_4H:C_6H = 1:0.053$ and $HC_5N:HC_7N:HC_9N = 1:0.16:0.25$ in L1544, and 1:0.048 and 1:0.14: <0.28 in IRAS 16293, respectively). When compared to other prestellar and protostellar sources at different evolutionary stages, cyanopolyynes and polyynyl radicals show the same exponentially decreasing trend, suggesting that their chemistry is likely the same across different cold environments.

(vi) We used a gas-phase astrochemical model to reproduce the abundances and investigate the formation of carbon chains. The model employing $A_v = 10$, standard $\zeta_{CR}$, and typical C/O ratio reproduces the c-$C_3H$ abundance in both sources, and the $C_6H$ abundance in L1544. Conversely, while $HC_5N$ is well reproduced in IRAS 16293, it is underestimated by more than an order of magnitude in L1544. The discrepancy is even more pronounced for $HC_7N$, whose abundance is systematically underpredicted by about 2 orders of magnitude in both sources.

(vii) Models reproduce the $C_6H/C_4H$ and $HC_5N/HC_3N$ ratios in both L1544 and IRAS 16293 for $A_v = 10$ and $t = 2–7 \times 10^5$ yr. However, the $HC_7N/HC_5N$ and $HC_9N/HC_7N$ ratios are systematically underestimated, likely due to incomplete chemical networks for large cyanopolyynes.

This work paves the way for future astrochemical studies with next-generation radio interferometers such as SKAO. By enabling high-angular resolution observations of complex carbon-bearing molecules, these facilities will allow us to probe chemical complexity down to 100 au scales within star-forming regions. Combined with updated chemical reaction networks, such observations will provide critical insights into the formation carbon chains and cyanopolyynes and their prebiotic role.

## ACKNOWLEDGEMENTS

We thank the anonymous referee for the instructive comments which improved the paper. The authors warmly acknowledge Angelique Kahle and Friedrich Wyrowski for providing the CO (3-2) APEX data. This project has received funding from the EC H2020 research and innovation programme for: (i) the project 'Astro-Chemical Origins' (ACO, No. 811312), and (ii) the European Research Council (ERC) project 'The Dawn of Organic Chemistry' (DOC, No. 741002). ClCo, LP, and GS acknowledge the PRIN-MUR 2020 BEYOND-2p (Astrochemistry beyond the second period

<footnote n="8">⁸www.sarao.za/science/meerkat/</footnote>

<footnote n="9">⁹www.skao.int/en</footnote>
elements, Prot. 2020AFB3FX), the project ASI-Astrobiologia 2023 MIGLIORA (Modeling Chemical Complexity, F83C23000800005), the INAF-GO 2024 fundings ICES, and the INAF-GO 2023 fundings PROTO-SKA (Exploiting ALMA data to study planet-forming discs: preparing the advent of SKA, C13C23000770005). LP, CICo, and GS also acknowledge financial support under the National Recovery and Resilience Plan (NRRP), Mission 4, Component 2, Investment 1.1, Call for tender No. 104 published on 2.2.2022 by the Italian Ministry of University and Research (MUR), funded by the European Union–NextGenerationEU-Project Title 2022JC2Y93 Chemical Origins: linking the fossil composition of the Solar system with the chemistry of protoplanetary discs – CUP J53D23001600006 – Grant Assignment Decree No. 962 adopted on 30.06.2023 by the Italian Ministry of Ministry of University and Research (MUR). GS also acknowledges support from the INAF-Minigrant 2023 TRIESTE (‘TRacing the chemical Heritage of our origins: from protostars toplanets’). EB acknowledges the support from the Italian Ministry for Universities and Research under the Italian Science Fund (FIS 2 Call–Ministerial Decree No. 1236 of 1 August 2023) and the Next Generation EU funds within the National Recovery and Resilience Plan (PNRR), Mission 4–Education and Research, Component 2–From Research to Business (M4C2), Investment Line 3.1–Strengthening and creation of Research Infrastructures, Project IR0000034–‘STILES – Strengthening the Italian Leadership in ELT and SKA’. The Green Bank Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. LG acknowledges support from the COST Action CA22133 PLANETS.

DATA AVAILABILITY

The data underlying this article are available in the NRAO archive at https://data.nrao.edu/portal/#/, and can be accessed with the project code.

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## APPENDIX A: INCOMPLETE ION-NEUTRAL CHEMISTRY IN CYANOPOLYNE NETWORKS

Astrochemical reaction networks are typically based on the two most widely used data bases, (Kinetic Database for Astrochemistry; V. Wakelam et al. 2024) and UDFA (UMIST Database for Astrochemistry; T. J. Millar et al. 2024). The GRETOPABE network adopted in this work is also KIDA-based, updated with recent results from our group and others on selected iCOMs (e.g. J.-C. Loison, V. Wakelam & K. M. Hickson 2014b; N. Balucani, C. Ceccarelli & V. Taquet 2015; D. Neufeld et al. 2015; F. Vazart et al. 2016, 2020; D. Skouteris et al. 2018; S. Blázquez et al. 2020; C. Codella et al. 2020; L. Giani et al. 2023). Nevertheless, for many complex species with limited experimental and theoretical studies, such as cyanopolyyynes, the networks remain incomplete, leading to large uncertainties in model predictions. For example, as one moves from $HC_3N$ to $HC_{11}N$, the number of known production reactions decreases drastically
(from 46 to 27). While the neutral-neutral chemistry of HC₅N has recently been revised (L. Giani et al. 2025), the ion-neutral pathways remain poorly constrained. Ionic reactions thought to be relevant for HCₙN formation, particularly at times <10⁵ yr, include the electron recombination processes HCₙNH⁺ + e⁻ and H₃CₙN⁺ + e⁻ (W. D. Geppert et al. 2004; J.-C. Loison et al. 2017), yet the networks of the corresponding protonated precursors are largely incomplete. For instance, current data bases list 28 formation reactions for HC₃NH⁺, 5 for HC₅NH⁺, and only 4 for HC₇NH⁺ and HC₉NH⁺. In the case of H₃CₙN⁺ ions, four reactions are listed for each species, but none have been studied theoretically or experimentally, and identical rates and products are assumed in each case. In our network, the formation reactions of HC₇NH⁺ and HC₉NH⁺ from HC₇N⁺ and HC₉N⁺ are missing, even though these are among the key reactions responsible for the formation of HC₃N and HC₅N. While the most recent updates of the KIDA and UDfA data bases now include these reactions, their rate coefficients remain uninvestigated. We therefore suggest revising

both the formation reactions of HCₙNH⁺, namely, HCₙN⁺ + H₂, and those leading to HCₙN⁺ itself, such as CₙH₃⁺ + N and CₙN⁺ + H₂. The omission of these reactions, as well as the use of uncertain rate coefficients, may explain why our models underestimate the abundances of the larger cyanopolyynes. In particular, we expect that including these missing pathways could reduce the discrepancy of about one to two orders of magnitude between the model predictions and the observed abundances of the longer cyanopolyynes. However, arbitrarily adding reactions without properly revising their products and rates would introduce even larger uncertainties. We therefore strongly recommend a systematic revision of the ionic chemistry of cyanopolyynes to make their formation networks more complete, a task we plan to address in future work.

This paper has been typeset from a TeX/LATEX file prepared by the author.