2. Early evolution of SN 2001ig

SN 2001ig was discovered by Evans, White, & Bembrick (Reference Evans, White and Bembrick2001)Footnote ^a on 2001 December 10.43 UT, in the face-on spiral galaxy NGC 7424. A likely explosion date is 2001 December 3 (MJD 52246) (Ryder et al. Reference Ryder, Sadler, Subrahmanyan, Weiler, Panagia and Stockdale2004), based on the early radio light curve evolution. Following Soria et al. (Reference Soria, Cheng, Pakull, Motch and Russell2024), we adopt a galaxy distance $d = 10.8$ Mpc (taken as the average between the Hubble Flow distance of 10.1 Mpc and the Tully-Fisher distance of 11.5 Mpc, as reported in the NASA/IPAC Extragalactic Database); this corresponds to a distance modulus of 30.17 mag and a scale of 52 pc arcsec $^{-1}$ . The SN position, determined from Australia Telescope Compact Array (ATCA) maps (Soria et al. Reference Soria, Cheng, Pakull, Motch and Russell2024), is R.A.(J2000) $= 22^{\rm h}\,57^{\rm m}\,30^{\rm s}.74 (\pm 0{\kern.5pt}.^{{\kern-2.5pt}\prime\prime}04)$ , Dec.(J2000) $= -41^{\circ}\,02^{\prime}\,26{\kern.5pt}.^{{\kern-2.5pt}\prime\prime}35 (\pm 0{\kern.5pt}.^{{\kern-2.5pt}\prime\prime}05)$ .

Figure 1.

9 GHz ATCA contours (cyan) of SN 2001ig observed on 2024 April 24, overlaid on a Gemini/GMOS colour composite image from 2004 September 13. Contours are defined as $2^{n/2}$ times the local rms noise level, which was $\approx$ 12 $\mu$ Jy. The three plotted contours correspond to 68 $\mu$ Jy (5.7 $\sigma$ ), 96 $\mu$ Jy (8 $\sigma$ ) and 135 $\mu$ Jy (11.3 $\sigma$ ). The beam size as shown at lower right was 1 ${\kern.5pt}.^{{\kern-2.5pt}\prime\prime}$ 6 $\times$ 0 ${\kern.5pt}.^{{\kern-2.5pt}\prime\prime}$ 8 at a position angle (north through east) of 158 ${\kern.5pt}.^{{\kern-2.5pt}\circ}$ 0. See Ryder et al. (Reference Ryder, Murrowood and Stathakis2006, Reference Ryder2018) for a discussion on the nature of the optical counterpart.

Table 1.

Summary of the post-2003 ATCA observations of SN 2001ig.

Tentative identification of 2001ig as a Type IIb was already suggested from optical spectra taken from the 6.5-m Baade Magellan Telescope at Las Campanas Observatory on 2001 December 11–13 (Phillips et al. Reference Phillips, Suntzeff, Krisciunas, Carlberg, Gladders, Barrientos, Matheson and Jha2001). The weakening of the H lines and the general similarity with the template Type Ib SN 1993J (but with faster ejecta velocity) was confirmed with spectra taken from the 3.6-m telescope at the European Southern Observatory over the following month (Clocchiatti & Prieto Reference Clocchiatti and Prieto2001; Clocchiatti Reference Clocchiatti2002). Intensive follow-up optical spectroscopic studies over the first couple of years, with Keck (Filippenko & Chornock Reference Filippenko and Chornock2002; Silverman et al. Reference Silverman, Mazzali, Chornock, Filippenko, Clocchiatti, Phillips, Ganeshalingam and Foley2009), Gemini-South (Ryder et al. Reference Ryder, Murrowood and Stathakis2006) and the Very Large Telescope (Maund et al. Reference Maund, Craig Wheeler, Patat, Wang, Baade and Höflich2007), provided comprehensive coverage of the early hydrogen-rich photospheric phase, then of the disappearance of the hydrogen lines, and finally (after $\approx$ 250 days from the explosion) of the nebular phase, dominated by emission lines such as [Mg I] $\lambda4571$ , [O I] $\lambda\lambda6300,6364$ , and [Ca II] $\lambda\lambda7291,7324$ (similar to typical Ib SNe). A Gemini/GMOS spectrum at an age of 6 yr (Ryder et al. Reference Ryder2018) showed narrow He II $\lambda4686$ emission, similar to that seen to emerge in SN 2014C 1–2 yr post-explosion (Milisavljevic et al. Reference Milisavljevic2015), and interpreted as continuing interaction with a dense CSM. Broad-band optical photometry, based on Gemini/GMOS images from 2004 September 13 (age $\approx$ 1 000 days), shows a yellow counterpart (Fig. 1) consistent with a late-B through late-F supergiant (Ryder et al. Reference Ryder, Murrowood and Stathakis2006). However, Hubble Space Telescope Wide Field Camera 3 images in the near-UV (F275W and F336W filters) from 2016 April 28 (age $\approx$ 14.4 yr) showed a fainter, bluer point-like source (Ryder et al. Reference Ryder2018), consistent with a B2 type ( $19\,000\,{\mathrm K} \lesssim T_\textrm{eff} \lesssim 22\,000\,{\mathrm K}$ ) main sequence star. The discrepancy is possibly due to residual emission from the shocked gas in the earlier observations, while the later observations are consistent with a $\approx$ 9 M $_\odot$ companion star. The likely presence of a surviving companion star supports the scenario that hydrogen stripping in the progenitor of SN 2001ig was caused by binary interaction.

In the radio bands, high cadence observations were taken over the first $\sim$ 3 yr with the ATCA at 1.4, 2.4, 4.8 and 8.6 GHz, with a few, sparser datapoints at 18.8 GHz (Ryder et al. Reference Ryder, Sadler, Subrahmanyan, Weiler, Panagia and Stockdale2004). In addition, the SN was observed a few times with the Very Large Array (VLA) at 1.4, 4.9, 8.5, 15 and 22.5 GHz (Ryder et al. Reference Ryder, Sadler, Subrahmanyan, Weiler, Panagia and Stockdale2004). The data showed (Ryder et al. Reference Ryder, Sadler, Subrahmanyan, Weiler, Panagia and Stockdale2004) a general early trend consistent with the model of Weiler et al. (Reference Weiler, Panagia, Montes and Sramek2002) and Sramek & Weiler (Reference Sramek, Weiler and Weiler2003), with an initial rise (optically thick phase) followed by an asymptotic power-law decline (optically thin phase), with $F_\nu \propto (t-t_0)^{-1.5}$ . In the optically thin phase, the spectral index settles at $\alpha \approx -1$ . However, there is an interesting modulation, with a period of $\approx 150$ d, superposed on the standard flux evolution (Ryder et al. Reference Ryder, Sadler, Subrahmanyan, Weiler, Panagia and Stockdale2004). This was interpreted as a density modulation of the CSM encountered by the expanding blast wave. One scenario is a series of shell ejections or increased mass loss at regular time intervals over the last few thousand years of the progenitor’s life. This may be the mechanism responsible for the dense dust shells seen around the WR-OB binary WR 112 (Lau et al. Reference Lau, Hankins, Schödel, Sanchez-Bermudez, Moffat and Ressler2017) and the recurrent (optical) re-brightenings over the first year of Type IIn supernova iPTF13z (Nyholm et al. Reference Nyholm2017). Alternatively, Ryder et al. (Reference Ryder, Sadler, Subrahmanyan, Weiler, Panagia and Stockdale2004) suggested that the density enhancements follow a spiral structure (pinwheel or lawn sprinkler model) typical of dust produced by wind-wind collisions near periastron in eccentric massive binary systems. The dusty pinwheel model was theoretically illustrated via hydrodynamical simulations for example by Schwarz & Pringle (Reference Schwarz and Pringle1996) and Walder & Folini (Reference Walder, Folini, van der Hucht, Herrero and Esteban2003), and successfully applied to several observed WR-OB systems (Tuthill, Monnier, & Danchi Reference Tuthill, Monnier and Danchi1999; Monnier, Tuthill, & Danchi Reference Monnier, Tuthill and Danchi1999; Marchenko et al. Reference Marchenko, Moffat, Vacca, Côté and Doyon2002; Soulain et al. Reference Soulain2018).

3. Results of the new radio observations

To investigate the behaviour of SN 2001ig at ages $\gt10$ yr, we used several sets of ATCA observations in the 6-km array configurationFootnote ^b (Table 1); some of the observations were already in the public archives, while others were specifically obtained for this project. These new data were taken between 2013 June and 2024 April, with 2.1 GHz data being taken throughout that entire span, and 5.5 and 9 GHz data being taken from 2021 April until 2024 April, where the 5.5 and 9 GHz data are recorded simultaneously. All new datasets were taken with a bandwidth of 2 GHz around each central frequency (with the Compact Array Broadband Backend; Wilson et al. Reference Wilson2011), compared with the 128 MHz bandwidth of the 2001–2003 data. Details and results from the new ATCA observations are shown in Table 1.

Flux density calibration was done with the primary ATCA calibrator PKS B1934 $-$ 638, and phase calibration with the nearby source PKS B2310 $-$ 417. We processed data following standard proceduresFootnote ^c within the Common Astronomy Software Application (casa, version 5.1.2; CASA Team et al. Reference Team2022). Imaging used the casa task clean, using a Briggs robust parameter of 0 (Briggs Reference Briggs1995), which balances sensitivity and resolution. With this choice, we obtained the synthesised beam shapes listed in Table 1.

To measure the flux density of SN 2001ig, we fit for a point source in the image plane. To do this, we used the casa task imfit to fit a 2D Gaussian with a full-width-at-half-maximum fixed to the parameters of the synthesised beam. All flux density errors include a systematic uncertainty on the absolute flux density of 2% (e.g. Massardi et al. Reference Massardi, Bonaldi, Bonavera, López-Caniego, de Zotti and Ekers2011; McConnell et al. Reference McConnell, Sadler, Murphy and Ekers2012), which was added in quadrature with the measured noise in a source-free region close to the source position.

Figure 2.

Flux density of SN 2001ig at 2.4, 4.85 and 8.55 GHz, based on ATCA observations, compared with the canonical evolution model of Weiler et al. (Reference Weiler, Panagia, Montes and Sramek2002). The ATCA data from 2001–2003 (Ryder et al. Reference Ryder, Sadler, Subrahmanyan, Weiler, Panagia and Stockdale2004) were taken at central frequencies of 2.4 GHz (red datapoints and dashed line model), 4.85 GHz (magenta datapoints and dash-dotted line model) and 8.55 GHz (blue datapoints and dotted line model), with a bandwidth of 128 MHz. The data from 2013, 2021 and 2024 were taken at central frequencies of 2.1, 5.5 and 9 GHz (Table 1) and then rescaled to 2.4, 4.85 and 8.55 GHz with an assumed spectral index $\alpha = -1$ , for the purposes of this plot. The bandwidth of the 2013, 2021 and 2024 measurements is 2 GHz.

We note that ATCA observations taken on 2024 February 13 suffered from significant phase decorrelation due to poor observing conditions (weather). To address this in the 5.5 GHz data, where there were a number of bright sources in the field, we carried out iterations of phase-only self-calibration, down to a solution interval of 60 s (shorter time intervals did not further improve the image). For the 9 GHz data, we were not able to confidently self-calibrate due to the lack of any sufficiently bright sources in the field (all field sources were $\lt$ 5 mJy at 9 GHz). Therefore, to estimate the effects of the phase decorrelation, we re-calibrated the data treating every second scan of the phase calibrator as a ‘target’ and every other phase calibrator scan as the ‘phase calibrator’. Doing so reveals the ‘target’ phase calibrator scans were fainter by a factor of 0.52, indicating strong phase decorrelation. We then correct for this by applying that correction factor to the initial measurement of SN 2001ig ( $1.0 \pm 0.1$ mJy).

We also detected SN 2001ig in a number of Australian SKA Pathfinder (ASKAP; Hotan et al. Reference Hotan2021) observations. These observations were conducted as part of the Rapid ASKAP Continuum Survey (RACS; Hale et al. Reference Hale2021) and Variables and Slow Transients (VAST; Murphy et al. Reference Murphy2021) projects and obtained from the CSIRO ASKAP Science Data ArchiveFootnote ^d (CASDA). Both surveys use the ASKAPSoft pipeline (Cornwell, Voronkov, & Humphreys Reference Cornwell, Voronkov, Humphreys, Bones, Fiddy and Millane2012) for data reduction and the Selavy algorithm (Whiting Reference Whiting2012) to produce fitted source catalogues. RACS and VAST observations, respectively, have $\sim$ 15 min and $\sim$ 12 min integrations, both with a bandwidth of 288 MHz. All epochs of VAST are observed at a central frequency of $887.5$ MHz. RACS includes epochs observed at central frequencies of $887.5, 943.5,$ and $ 1367.5$ MHz. The ASKAP detections are shown in Fig. 3. Flux density errors are calculated as the quadrature sum of the fitted error, RMS, and a 6% flux scale uncertainty.

The main results of the late-time radio observations can be summarised as follows:

1. (i) there is strong re-brightening compared with the extrapolated power-law decline of 2001–2002. Today’s radio flux density is two orders of magnitude higher than we expect assuming a simple power-law model. Low-frequency (2.1 GHz) measurements show (Fig. 2) that the re-brightening must have started already before 2013 (age of $\approx$ 4 000 d) and more likely, as early as at an age of $\approx$ 1 000 d. At the assumed distance of 10.8 Mpc, the luminosity density at 5.5 GHz was $\approx$ 4.6 $\times 10^{26}$ erg s $^{-1}$ Hz $^{-1}$ in 2004 April.

2. (ii) the re-brightening is still ongoing. There is a slight but significant increase over the past $\approx$ 2 000 d, seen both from ATCA (Fig. 2 and Table 1) and from ASKAP (Fig. 3).

3. (iii) the spectral index of the optically thin synchrotron spectrum has remained settled around the canonical value $\alpha \approx -1$ , a value it had already attained after $\approx$ 1 yr (Ryder et al. Reference Ryder, Sadler, Subrahmanyan, Weiler, Panagia and Stockdale2004). Specifically, the ATCA and ASKAP observations from 2024 April suggest a spectral index $\alpha \approx -0.93 \pm 0.05$ between 0.89 GHz and 9.0 GHz, $\alpha \approx -0.97 \pm 0.05$ between 2.1 GHz and 9.0 GHz, and $\alpha \approx -1.07 \pm 0.05$ between 5.5 GHz and 9.0 GHz. Thus, the re-brightening did not correspond to a change in the spectral slope.

Figure 3.

ASKAP detections of SN 2001ig at late times, starting around $\approx 6\,300$ days. The points are coloured by their central observing frequency, with the majority of the observations coming from the $887.5$ MHz VAST survey.

The first hint of an increased CSM interaction leading to a re-brightening was already noted by Ryder et al. (Reference Ryder2018), from the stack of two Gemini/GMOS spectra taken on 2007 July 18 and November 6 (average age $\approx$ 2 100 d). The combined spectrum showed narrow HeII $\lambda4686$ emission, a feature sometimes seen in re-brightening SNe.

4. Modelling

We assume that the radio emission is made of an optically thick and an optically thin component (Chevalier Reference Chevalier1998). In the optically thick limit, the flux density, $S_{\nu}$ , is proportional to the observing frequency $\nu$ such that $S_{\nu} \propto \nu^{5/2}$ , while in the optically thin limit, $S_{\nu} \propto \nu^{-(p-1)/2}$ , where p, the spectral index of the particle energies, is $\gt 2$ (typically, $p \sim 3$ ). The peak flux density at a given time is a function of the shock radius $R_{s}$ and the magnetic field strength B (Chevalier Reference Chevalier1998). From $R_{ s}(t)$ , we can infer the shock velocity $v_{s}$ , the CSM density $\rho_\textrm{csm}$ and other quantities of interest for the evolution of the system.

We model SN 2001ig in two slightly different ways. First, as an initial approximation, we take the observed peak flux density at a given frequency, and assume a value of $p \equiv 3$ and a constant shock velocity to derive $\rho_\textrm{ csm}$ . The main purpose of this model is to estimate $\dot{M}/v_{w}$ of the progenitor star. Then, we use a Markov Chain Monte Carlo (MCMC) model to fit $\rho_\textrm{csm}$ at each epoch in order to produce the observed flux density. The main purpose of this model is to estimate the density enhancement encountered by the shock wave at late times.

4.1 Determining the mass loss rate of the progenitor

We assume energy equipartition between the magnetic field and the relativistic electrons ( $\alpha = 1$ in the formalism of Chevalier Reference Chevalier1998), with a 10% fraction of post-shock energy going into each component: $\epsilon_B = \epsilon_e = 0.1$ . We also assume $p \equiv 3$ , corresponding to an optically thin synchrotron spectral index $\alpha = -1$ , and an emission filling factor (fraction of a spherical volume with outer radius $R_{s}$ ) $f = 0.5$ . Then, the outer shock radius and the magnetic field are obtained from Eqs. (13–14) of Chevalier (Reference Chevalier1998), as a function of the peak flux density $S_{\nu,{p}}$ , namely

(1) \begin{align}R_{s,p} & = 8.8 \times 10^{15} \left(\frac{\epsilon_B}{\epsilon_e}\right)^{1/19}\left(\frac{f}{0.5}\right)^{-1/19} \left(\frac{S_{\nu,{p}}}{{\mathrm {Jy}}}\right)^{9/19}\nonumber \\& \left(\frac{d}{{\mathrm {Mpc}}}\right)^{18/19}\left(\frac{\nu}{5{\mathrm {\,GHz}}}\right)^{-1}\ {\mathrm {{cm}}}\end{align}

(2) \begin{align}B_{p} & = 0.58 \left(\frac{\epsilon_B}{\epsilon_e}\right)^{4/19}\,\left(\frac{f}{0.5}\right)^{-4/19}\, \left(\frac{S_{\nu,{p}}}{{\mathrm {Jy}}}\right)^{-2/19} \nonumber\\&\left(\frac{d}{{\mathrm {Mpc}}}\right)^{-4/19}\,\left(\frac{\nu}{5{\mathrm {\,GHz}}}\right)^{-1}\ {\mathrm G}\end{align}

For $S_{\nu,{p}}$ ,we take the brightest flux density measurement available at any frequency: $S_{18.8} \approx 43$ mJy beam $^{-1}$ , from the ATCA at 18.8 GHz at an age $t-t_0 \approx 28$ d. Thus, the radius of the shocked bubble at that time is $R_{s,p} \approx 5.0 \times 10^{15}$ cm, with a field $B_{p}\approx 1.8$ G. The shock velocity is then simply $v_{s} = R_{s,p}/(t-t_0) \approx 20\,000$ km s $^{-1}$ . To derive the mass loss rate from the progenitor, we use the standard assumption (adopted for example in Rose et al. Reference Rose2024) that the magnetic energy density $B^2/8\pi$ is a fraction $\epsilon_B$ of the post shock energy density: $B^2/8\pi \approx (9/8)\,\epsilon_B \rho_\textrm{csm}v_{s}^2$ . From Eq. (19) of Chevalier (Reference Chevalier1998), for a constant mass-loss rate in a steady progenitor wind and constant shock velocity, we have:

(3) \begin{eqnarray}\dot{M} & \approx & 5.2 \times 10^{-6}\, \left( \frac{\epsilon_B}{0.1}\right)^{-1} \left( \frac{B_{p}^2}{1\ {\mathrm G}} \right)^{2}\left( \frac{t-t_0}{10\ {\mathrm d}}\right)^{2}\nonumber \\&&\left( \frac{v_{w}}{1\,000\ {\mathrm {km~s}}^{-1}}\right)\ \ M_\odot {\rm yr}^{-1}\nonumber \\&\approx& 1.3 \times 10^{-4}\left( \frac{v_{w}}{1\,000\ {\mathrm {km~s}}^{-1}}\right)\ \ {\rm M}_\odot {\rm yr}^{-1}.\end{eqnarray}

At the peak flux density, the CSM density is

(4) \begin{equation}\rho_\textrm{csm} \equiv \frac{\dot{M}}{4\pi R^2_{s,p}v_{w}} \approx 3 \times 10^{-19}\ \ {\mathrm {g~cm}}^{-3}.\end{equation}

Comparing these values with the physical quantities derived in the ASKAP survey of Rose et al. (Reference Rose2024), we note that the Type Ic SN 2016coi (Grupe et al. Reference Grupe, Brown, Dong, Shappee, Holoien, Stanek, Prieto and Margutti2016; Terreran et al. Reference Terreran2019) appears the most similar to 2001ig at early times.

4.2 Determining the late-time density enhancement

To model the evolution of the synchrotron emission from early to late times, we adopt the thin shell approximation (Cox Reference Cox1972; Castor, McCray, & Weaver Reference Castor, McCray and Weaver1975; Weaver et al. Reference Weaver, McCray, Castor, Shapiro and Moore1977), and solve, via numerical integration, a simultaneous system of equations for the shock radius and velocity at each epoch, the magnetic field, and the peak flux of the synchrotron radiation that results from shock interactions. Specifically, the thin shell differential equation is

(5) \begin{equation} M_{s} \frac{dv_{s}}{dt} = 4\pi R_{s}^2[\rho_{\textrm{ej}}(v_{\textrm{ej}} - v_{s})^2 - \rho_{\textrm{csm}}v_{s}^2]\end{equation}

(Chevalier & Fransson Reference Chevalier, Fransson, Alsabti and Murdin2017), where $M_{s}$ is the mass of the thin shell comprising shocked ejecta and shocked CSM at the shock radius $R_{s}$ , propagating at the shock velocity $v_{s} = dR_{s}/dt$ ; $v_{\textrm{ej}}$ is the ejecta velocity at the reverse shock; $\rho_{\textrm{ej}}$ and $\rho_{\textrm{csm}}$ are the density profiles of the shocked ejecta and CSM, respectively. $M_{s}$ is computed from the numerical integration of the density profiles. For the CSM density, we assume again that the post-shock energy density is proportional to the magnetic energy density ( $\rho_\textrm{csm}v_{s}^2 \propto B^2/\epsilon_B$ as in Section 4.2). For the ejecta density profile we adopt

(6) \begin{equation}\rho_{\textrm{ej}} =\left\{ \begin{array}{lr} F(t-t_0)^{-3}, & v_{\textrm{ej}} \lt v_t\\ F(t-t_0)^{-3}\left(v_{\textrm{ej}}/v_t\right)^{-n}, & v_{\textrm{ej}} \geq v_t \end{array}\right.\end{equation}

where

(7) \begin{equation} F = \frac{1}{4\pi n} \frac{[3(n-3)M_{\textrm{ej}}]^{5/2}}{[10(n-2)E]^{3/2}},\end{equation}

(8) \begin{equation} v_t = \biggl[\frac{10(n-5)E}{3(n-3)M_{\textrm{ej}}}\biggl]^{1/2},\end{equation}

and $M_\textrm{ej}$ and E are the total explosion mass and energy. Given the uncertainty on the total ejecta mass, we will solve for the CSM profile for a low end estimate of $M_{ej} = 1\,{\rm M}_\odot$ (approximately the ejecta mass estimated by Silverman et al. Reference Silverman, Mazzali, Chornock, Filippenko, Clocchiatti, Phillips, Ganeshalingam and Foley2009) and a high end estimate of $M_{ej} = 4\,{\rm M}_\odot$ ; stripped-envelope SN models by Aguilera-Dena et al. (Reference Aguilera-Dena, Müller, Antoniadis, Langer, Dessart, Vigna-Gómez and Yoon2023) suggest that the mass range between 1 and 4 M $_\odot$ covers about 80% of events. For the 1 M $_\odot$ model, we assume a fixed explosion energy of $E = 10^{51}$ erg, corresponding to an initial ejecta velocity of $v_i \approx 10\,000$ km/s. For the 4 M $_\odot$ model, instead, we find that the using the same explosion energy results in an initial ejecta velocity, $v_i \approx 5\,000$ km/s that is too slow to reproduce the observed early time emission with any density profile. In order to obtain an acceptable flux density fit for this high value of $M_{\textrm{ej}}$ , we allow the initial ejecta velocity to be a free parameter, $v_{{i}}$ . For this higher mass scenario we find that an initial ejecta velocity of $v_i = 13\,888$ km/s is needed to replicate the early time emission. Finally, we adopt an ejecta density index $n = 10$ , suitable for massive stars with radiative envelopes and convective cores (Chevalier Reference Chevalier1982a; Matzner & McKee Reference Matzner and McKee1999; Chevalier & Fransson Reference Chevalier, Fransson, Alsabti and Murdin2017), including Wolf-Rayet stars (Kippenhahn & Weigert Reference Kippenhahn and Weigert1990).

In addition to Eq. (5) for the shock dynamics, and the equations for ejecta and CSM density, we simultaneously solve Eqs. (11–12) from Chevalier (Reference Chevalier1998) to infer the peak flux density $S_{\nu,{p}}$ and peak frequency $\nu_{p}$ from the calculated radius and magnetic field strength at each epoch. As we did in Section 4.1, we assume that the fractions of post-shock energy transferred to the magnetic field and the relativistic electrons are $\epsilon_B = 0.1$ and $\epsilon_e = 0.1$ , respectively, and that these fractions do not evolve with time; we assume again a standard filling factor, $f = 0.5$ . We then model the radio spectrum as a broken power law, with an optically thick slope of 5/2 and an optically thin slope of $(1-p)/2$ :

(9) \begin{equation} S_\nu = 2 S_{\nu,{p}} \left[ \left(\frac{\nu}{\nu_{{p}}}\right)^{-(1-p)/2} + \left(\frac{\nu}{\nu_{{p}}}\right)^{-5/2}\right]^{-1}.\end{equation}

Unlike what we did in Section 4.1, here we do not assume a spectral index $p=3$ for the electron energies; instead, we fit our broken power-law spectrum across all epochs with an MCMC fit (with the python package emcee Footnote ^e ), and find that the best fitting value is constrained to be $p = 2.83 \pm 0.02$ .

We assume that at early times, $(t-t_0) \lt 700$ d, the SN shock is interacting with a CSM density profile that scales as $\rho_{\textrm{csm}}\propto R^{-\alpha}$ . Using an MCMC to model the shock dynamics and resulting radiation for just the early time data, we infer $\alpha = 1.96\pm 0.01$ . We thus assume a standard wind density profile ( $\rho_{\textrm{csm}} \propto R^{-2}$ ) at smaller radii. Our main objective here is to explain the observed radio re-brightening at late times (Fig. 2). We want to test whether a deviation from the CSM wind-density profile at large radii can explain the re-brightening, and if so, what amount of density enhancement may be needed. To this aim, we introduce a break radius $R_{\rm{brk}}$ , where the shock encounters an over-density compared to the extrapolation of the wind-density profile, and starts to interact with a CSM characterised by shallower density index, $\alpha_{\rm{over}}$ . For a detailed description of the model, in the more general case of a three-zone CSM, see Ibik et al. (Reference Ibik2024).

We use another MCMC fit to solve for the CSM profile that would produce the observed emission. The fit parameters are: $R_{\textrm{brk}}$ , $\rho_{\textrm{0,wind}}$ , $\rho_{\textrm{0,over}}$ , and $\alpha_{\textrm{over}}$ . $R_{\textrm{brk}}$ represents the radius at which the shock encounters an overdensity; $\rho_{\textrm{0,wind}}$ and $\rho_{\textrm{0,over}}$ are the CSM densities immediately before and after $R_{\textrm{brk}}$ . We assume that the CSM follows a wind density profile ( $\rho_{\textrm{csm}} \propto R^{-2}$ ) at $R \lt R_{\textrm{brk}}$ and a shallower profile $\rho_{\textrm{csm}} \propto R^{-\alpha_{\textrm{over}}}$ at $R \gt R_{\textrm{brk}}$ :

(10) \begin{equation}\rho_{\textrm{CSM}} =\left\{ \begin{array}{lr} \rho_\textrm{0, wind} \left(R/R_\textrm{brk}\right)^{-2}, & R \lt R_\textrm{brk}\\ \rho_\textrm{0, over} \left(R/R_\textrm{brk}\right)^{-\alpha_\textrm{over}}, & R \gt R_\textrm{brk} \end{array}\right.\end{equation}

For this treatment of the shock dynamics, we expect the shock velocity to decrease steadily over time, as the shock propagates through power law CSM profiles. At the density discontinuity $R_\textrm{brk}$ , the rate of deceleration will change as the shock encounters an overdensity, with a potentially different power law index, $\alpha_\textrm{over}$ . With this evolution, we expect the overall flux density to steadily decline as the shock velocity decreases along the wind density section of the CSM profile. When encountering the overdensity at $R_\textrm{brk}$ , the flux density will increase as the shock accelerates a larger population of electrons to produce synchrotron emission. At very late times ( $(t-t_0) \gtrsim 2 \times 10^{4}$ d and $(t-t_0) \gtrsim 5 \times 10^{4}$ d for the low and high $M_{\textrm{ej}}$ scenarios, respectively), the flux density will once again begin to decline, as the shock velocity keeps decreasing.

The best-fitting CSM density profiles for the low and high $M_{\textrm{ej}}$ scenarios are shown in Fig. 4, and the corresponding time evolutions of the model flux densities are plotted in Fig. 5 and (for the representative 5-GHz case) Fig. 6: lightcurve. Details of the fit results for each parameter are in the Appendix, Figs. A1–A2. Although the density jump at the break radius is somewhat small in both cases, this increase, in combination with the shallower power-law index $\alpha_{\rm{over}}$ , creates an overdensity of about an order of magnitude at the most recent epochs of observation (year 2024) compared to what it would have been in case of a standard wind profile. This, in turn, leads to radio flux densities about two orders of magnitude higher than expected from the extrapolation of the early-time decline. The best fit value for $\alpha_{\rm{over}}$ remains constant across both fits, at $\alpha_{\rm{over}} = 0.75$ . This shallow index, corresponding to a slow evolution of CSM density, is needed to reproduce the slow evolution of synchrotron emission after the re-brightening.

Figure 4.

Inferred CSM density profiles as a function of shock radius, for (a) $M_{\rm{ej}} = 1\,{\rm M}_\odot$ , and (b) $M_{\rm{ej}} = 4\,{\rm M}_\odot$ . In each panel, the blue line is our model density profile that provides the best fit to the observed radio datapoints at early and late times; the dashed red line is the density profile assuming a uniform stellar wind ( $\rho_{\textrm{csm}} \propto R^{-2}$ ). The green markers on the X axis represent the best-fitting shock radius at each of the epochs with radio observations. The gap in the radio coverage between ages of $\approx$ 2–12 yr creates an uncertainty about where the density enhancement occurs (i.e. the parameter $R_{\textrm{brk}}$ in our model, Section 4.2). The portions of the CSM density profiles well constrained by the radio observations are plotted as solid blue segments, while the intervals without a strong constraint are dashed.

Figure 5.

Spectral energy distributions of the best-fitting models for (a) $M_{\rm{ej}} = 1\,{\rm M}_\odot$ , and (b) $M_{\rm{ej}} = 4\,{\rm M}_\odot$ , at each time step. Datapoints (solid circles) are the observed radio flux densities, colour-coded by SN age, from red (earliest) to violet (latest). Consecutive observations with $\Delta t/t \lt 0.1$ are combined, and their plotted colour corresponds to the average age. Solid lines are model radio spectra (also colour-coded by age) corresponding to our best-fitting CSM density profile and shock velocity. At earlier times, flux density and peak frequency decrease over time, as expected from a slowly decelerating shock propagating through a uniformly expanding wind. At later times, the flux increases again, which we interpret as a flattening of the CSM density profile.

Figure 6.

Flux density evolution of the best-fitting spectral model for (a) $M_{\rm{ej}} = 1\,{\rm M}_\odot$ , and (b) $M_{\rm{ej}} = 4\,{\rm M}_\odot$ , for the representative 5-GHz case, compared with the observed datapoints (red circles).

As a caveat, we stress that the lack of data between $\approx$ 3–10 yr post-explosion creates uncertainty as to where exactly the overdensity begins. This can clearly be seen from the large gap between the fitted values of the shock radii at each epoch of observations, marked as green segments on the X axis of Fig. 4; the overdensity begins somewhere within that gap. As a result, although our best-fitting parameters represent one solution in good agreement with the observations, there are still strong degeneracies between $R_{\textrm{brk}}$ and the normalisations of $\rho_{\textrm{0,wind}}$ and $\rho_{\textrm{0,over}}$ . A different combination of break radius and overdensity normalisation could produce the same densities at the points where we actually have measurements.

The best-fitting density profiles in Fig. 4 reach fairly low densities (number densities $n_e \sim 10$ –100 cm $^{-3}$ ), especially in the case of high $M_{\textrm{ej}}$ . However, this is a direct result of our choice of the microphysical parameter $\epsilon_B = 0.1$ . Some models of non-relativistic shock acceleration predict values of $\epsilon_B$ an order of magnitude lower (Gupta, Caprioli, & Spitkovsky Reference Gupta, Caprioli and Spitkovsky2024). If we were to assume a lower fraction of the total energy transferred to the magnetic field strength, we would in turn infer higher densities, to produce the same observed radio emission. Therefore, our density models can best be interpreted as lower limits on the CSM density profile. While the density normalisation remains uncertain, the inferred trend of a flattening density index at late times is a solid result regardless of the choices of $M_{\textrm{ej}}$ and $\epsilon_B$ .