MNRAS 529, L33–L40 (2024) https://doi.org/10.1093/mnrasl/slad195 Advance Access publication 2023 December 28 The metamorphosis of the Type Ib SN 2019yvr: late-time interaction Luc ´ıa Ferrari , 1 , 2 ‹ Gast ´on Folatelli, 1 , 2 , 3 Hanindyo Kuncarayakti, 4 , 5 Maximilian Stritzinger, 6 Keiichi Maeda , 7 Melina Bersten, 1 , 2 , 3 Lili M. Rom ´an Aguilar, 1 , 2 M. Manuela S ´aez, 8 , 9 Luc Dessart, 10 Peter Lundqvist , 11 Paolo Mazzali, 12 , 13 Takashi Nagao , 4 , 14 , 15 Chris Ashall, 16 Subhash Bose , 6 , 17 Se ´an J. Brennan, 11 Yongzhi Cai, 18 , 19 , 20 Rasmus Handberg , 6 Simon Holmbo, 6 Emir Karamehmetoglu, 6 Andrea Pastorello, 21 Andrea Reguitti , 21 , 22 Joseph Anderson, 23 Ting-Wan Chen, 24 Llu ´ıs Galbany, 25 , 26 Mariusz Gromadzki , 27 Claudia P. Guti ´errez , 25 , 26 Cosimo Inserra , 28 Erkki Kankare, 4 Tom ´as E. M ¨uller Bra v o, 25 , 26 Seppo Mattila, 4 , 29 Matt Nicholl, 30 Giuliano Pignata, 31 Jesper Sollerman , 11 Shubham Sri v asta v 32 and Da vid R. Young 32 Affiliations are listed at the end of the paper Accepted 2023 December 12. Received 2023 December 7; in original form 2023 September 30 A B S T R A C T We present observational evidence of late-time interaction between the ejecta of the hydrogen-poor Type Ib supernova (SN) 2019yvr and hydrogen-rich circumstellar material (CSM), similar to the Type Ib SN 2014C. A narrow H α emission line appears simultaneously with a break in the light-curve decline rate at around 80–100 d after explosion. From the interaction delay and the ejecta velocity, under the assumption that the CSM is detached from the progenitor, we estimate the CSM inner radius to be located at ∼6.5–9.1 × 10 15 cm. The H α emission line persists throughout the nebular phase at least up to + 420 d post-explosion, with a full width at half maximum of ∼2000 km s −1 . Assuming a steady mass-loss, the estimated mass-loss rate from the luminosity of the H α line is ∼3–7 × 10 −5 M  yr −1 . From hydrodynamical modelling and analysis of the nebular spectra, we find a progenitor He-core mass of 3–4 M , which would imply an initial mass of 13–15 M . Our result supports the case of a relati vely lo w-mass progenitor possibly in a binary system as opposed to a higher mass single star undergoing a luminous blue variable phase. K ey words: supernov ae: general – supernov ae: indi vidual: SN 2019yvr. 1 S U o f U c ( t H o e H m o t ( S e h  2 a S ( e t F C e ( i o fi I a t o i T © P C p D ow nloaded from https://academ ic.oup.com /m nrasl/article/529/1/L33/7503375 by Turku U niversity - IN AC TIVE user on 18 January 2024 I N T RO D U C T I O N upernovae (SNe) are among the most powerful explosions in the niverse, and their study provides valuable insights into a multitude f astrophysical processes including stellar evolution, the subsequent ormation of compact objects, and the chemical enrichment of the niverse. A substantial fraction of SNe are associated with the ollapse of the iron cores of massive stars, called core-collapse SNe CCSNe). CCSNe are classified based on their spectral characteris- ics. Objects with H features are classified as Type II, those lacking but exhibiting He are of Type Ib, and those with no H nor He are f Type Ic (for a contemporary and concise re vie w, see Stritzinger t al. 2023b , and references therein). Although relatively rare, the -poor stripped-envelope SNe are of particular interest as they are ost lik ely link ed to the death of a He or C + O star, as an analogue f Wolf–Rayet stars, which experienced significant mass-loss either hrough strong stellar winds or interaction with a binary companion e.g. W oosley, Langer & W eaver 1993 ; Podsiadlowski et al. 2004 ; mith 2014 ). Moreo v er, it has been proposed that the progenitor nvelope could be removed by a combination of both processes, i. e. ybrid mass-loss (Fang et al. 2019 ; Sun, Maund & Crowther 2023 ). E-mail: luciaferrari4@gmail.com w s The Author(s) 2023. ublished by Oxford University Press on behalf of Royal Astronomical Society. Th ommons Attribution License ( http:// creativecommons.org/ licenses/ by/ 4.0/ ), whic rovided the original work is properly cited. A small but growing number of Type Ib (e.g. Milisavljevic et al. 015 ; Vinko et al. 2017 ; Mauerhan et al. 2018 ; Chandra et al. 2020 ) nd Type Ic SNe (Kuncarayakti et al. 2018 ; Tartaglia et al. 2021 ; tritzinger et al. 2023a ) exhibit signatures of circumstellar interaction CSI) in optical wavelengths. CSI occurs when rapidly expanding SN jecta shock a dense circumstellar material (CSM), originated from he progenitor star itself, or from a companion star (e.g. Che v alier & ransson 1994 ; Fransson et al. 2002 ; Yoon 2017 ). Signatures of SI in some cases are revealed in the spectra as narrow Balmer mission lines, reminiscent to the hallmark feature of Type IIn SNe e.g. Schlegel 1990 ; Taddia et al. 2013 ), as well as in some cases high- onization coronal lines, and/or excesses of flux in different regions f the electromagnetic spectrum (e.g. Stritzinger et al. 2012 ). In this paper, we examine SN 2019yvr, which was initially classi- ed as a Type Ib SN (Dimitriadis 2019 ), but e ventually de veloped SN In-like features. Based on pre-explosion Hubble Space Telescope rchi v al images, Kilpatrick et al. ( 2021 ) found a point source at he location of SN 2019yvr. The spectral energy distribution (SED) f the source suggests a cool and luminous progenitor candidate, n contradiction with the He star picture of a SN Ib progenitor. hey also discuss a binary scenario but also find it incompatible ith a Type Ib SN progenitor. Sun et al. ( 2022 ) explored several cenarios by performing an environmental study, and proposed a is is an Open Access article distributed under the terms of the Creative h permits unrestricted reuse, distribution, and reproduction in any medium, L34 L. Ferrari et al. M b a p i i p c ( p 2 S L a M T a N e s t m i f b v e C ( u b d u p t p m o o f t T S t E S o A s e s 1 2 3 4 Figure 1. SN 2019yvr transitions from a SN Ib spectrum around maximum to a Type IIn-like spectrum at late phase. Comparison data include a similar phase spectrum of the SN Ib iPTF13bvn (Sri v astav, Anupama & Sahu 2014 ) and the Type IIn SN 2006jd (Stritzinger et al. 2012 ). Spectra are corrected by redshift and not corrected by extinction. V 2 c S 5 t V o R r M 3 3 S 2 b M o 2 t s S t s T a b H a b e s D ow nloaded from https://academ ic.oup.com /m nrasl/article/529/1/L33/7503375 by Turku U niversity - IN AC TIVE user on 18 January 2024inary system, composed of a hot and compact SN progenitor nd a yellow hypergiant (YHG) companion. SN 2019yvr therefore rovides an opportunity to understand better the pre-SN evolution of nteracting events’ progenitors. The present work focuses on the appearance of the CSI features n SN 2019yvr, the late-time interaction, and particular progenitor roperties such as the pre-SN mass and mass-loss rate. We will omplement this analysis with an upcoming paper by Ferrari et al. in preparation) where we will provide a detailed analysis of the full hotometric and spectroscopic evolution. OBSERVATION S N 2019yvr was first reported by the Asteroid Terrestrial-impact ast Alert System on 2019 December 27.5 UT (Tonry et al. 2019 ), nd classified two days later as a Type Ib SN (Dimitriadis et al. 2019 ; uller et al. 2019 ). Based on the last non-detection from the Zwicky ransient Facility (Masci et al. 2019 ), we adopt an explosion epoch s JD = 2458839.89 ±3.84. The SN is located in the nearby Galaxy GC 4666, which hosted the Type Ia SN ASASSN-14lp (Shappee t al. 2016 ). We adopt their derived distance of 14.7 ± 1.5 Mpc. See ection 1 of the Supplementary Material for details on the disco v ery, he estimated explosion date and the adopted distance. The data employed in this study are part of a larger set of ultiband optical light curves and a series of spectra (Ferrari et al., n preparation). Here we present BVgri -band photometry measured rom images obtained by the NUTS 1 (Nordic optical telescope Un- iased Transient Surv e y) and ASAS-SN 2 (All-Sky Automated Sur- e y for Supernovae) collaborations using the 2.56-m NOT telescope quipped with ALFOSC (Andalucia Faint Object Spectrograph and amera) and the 1-m Las Cumbres Observatory Global Telescope LCOGT) network respectively. The NOT images were reduced sing the PYRAF -based ALFOSCGUI 3 reduction pipeline developed y E. Cappellaro, while fully processed LCOGT images were ownloaded from the observatory’s data archive. Host-galaxy subtraction was performed on all science frames sing template images obtained prior to 2019. Point-spread-function hotometry of SN 2019yvr was then measured relative to stars in he field using the Aarhus-Barcelona FLOWS project’s automated ipeline. 4 The photometry is tabulated in section 2 of the Supple- entary Material. We also present eight low-resolution and one high-resolution ptical spectra, which are summarized in the journal of spectroscopic bservations (section 2 of the Supplementary Material). In summary, our spectra were obtained by ePESSTO + (Smartt et al. 2015 ) with he European Southern Observatory (ESO) 3.58-m New Technology elescope (Buzzoni et al. 1984 ) equipped with the ESO Faint Object pectrograph and Camera optical (EFOSC2), three by NUTS with he NOT ( + ALFOSC), and one spectrum was obtained with the SO 8.4-m Very Large Telescope FOcal Reducer and low-dispersion pectrograph (Appenzeller et al. 1998 ), as part of the FORS + Surv e y f Supernovae in Late Times program (see Kuncarayakti et al. 2022 ). high-resolution spectrum was obtained with the 8.2-m Subaru Tele- cope equipped with the High Dispersion Spectrograph (Noguchi t al. 2002 ). The spectroscopic observations were reduced following tandard techniques using the respective instrument pipelines. NRASL 529, L33–L40 (2024) http:// nuts.sn.ie/ . https://www .astronomy .ohio-state.edu/asassn/. http:// graspa.oapd.inaf.it/ foscgui.html . https:// github.com/ SNflows . p 5 fl mThe Milky Way reddening along the line of sight is E ( B − ) MW = 0.022 mag (i.e. A MW V = 0 . 068 mag, Schlafly & Finkbeiner 011 , assuming R V = 3.1). After comparing the observed colour urves of SN 2019yvr with the intrinsic colour-curve templates from tritzinger et al. ( 2018 ), analysing the diffuse interstellar band at 780 Å and the Na I D lines in the high-resolution spectrum from he Subaru telescope, we adopt a host-galaxy reddening of E ( B − ) host = 0.57 ± 0.09 mag. This value is in agreement with the estimate f Kilpatrick et al. ( 2021 ) of 0 . 51 + 0 . 27 −0 . 16 mag and with that published in odr ´ıguez, Maoz & Nakar ( 2022 ) of 0.56 ± 0.09 mag. Details on the eddening estimation are presented in section 3 of the Supplementary aterial. SI GNATURES O F CSM I N T E R AC T I O N .1 Emergence of H α emission N 2019yvr showed H α emission at late times similarly to SN 014C, which in that case was interpreted as a result of interaction etween the ejecta and H-rich CSM (e.g. Milisavljevic et al. 2015 ; argutti et al. 2017 ). In Fig. 1 , we compare the spectra of SN 2019yvr btained + 2 d and + 383 d past the epoch of B -band maximum, JD = 458851.6 (see table 1 of the Supplementary Material). Throughout he paper we will refer the phases to this date unless otherwise pecified. While the first spectrum is consistent with that of typical Ne Ib, the late-phase spectrum clearly exhibits the transformation o a SN IIn-like spectrum. The temporal co v erage of SN 2019yvr observ ations allo ws us to tudy the moment when signatures of interaction become apparent. his is depicted in Fig. 2 where we show the evolution of the spectra round the He I λ6678 (which includes H α) and the He I λ7065 lines etween + 42 d and + 118 d. The + 118 d spectrum clearly exhibits α in emission, whereas the previous spectra show an absorption t the same wavelength. While the absorption due to He I λ6678 ecomes substantially weaker with time before the H α emission merges, the absorption due to He I λ7065 remains roughly constant. This phenomenon can be appreciated in the evolution of the ab- orption pseudo-equi v alent width (pEW) 5 of both lines. We measured EW using the splot task from IRAF five times to account for values The term ‘pseudo-equi v alent width’ is used because the actual continuum ux is unknown and we therefore fit a ‘pseudo-continuum’ between the local axima at both sides of the line. SN 2019yvr L35 Figure 2. Evolution of the He I λ6678 and He I λ7065 features of SN 2019yvr from + 42 d to + 118 d plotted in velocity space. Dashed vertical lines indicate the 0 km s −1 position, while in the left panel the dotted vertical line indicates the rest wavelength of H α. Complete spectra are displayed in section 4 of the Supplementary Material. Figure 3. Light curves of SN 2019yvr compared with those of SN 2011dh (Ergon et al. 2015 ) and iPTF13bvn (Fremling et al. 2016 ), shifted in magnitudes to coincide at the peak. Dashed lines correspond to 56 Co decay. The post-maximum light curves of SN 2019yvr begin to deviate from the normal decay line of the comparison objects beginning around + 80 d. w 4 r o b t i a s t d 3 T T o d i a p 2 S o i S i i i c b a 3 B t + w d m t 6 t i e i u e t w d t t t H t s 7 8 a t L 1 A m t fl c ∼ t e b s r D ow nloaded from https://academ ic.oup.com /m nrasl/article/529/1/L33/7503375 by Turku U niversity - IN AC TIVE user on 18 January 2024ith their corresponding errors. For He I λ6678, we obtain pEW = 0.4 ± 0.9, 26.9 ± 0.6, and 10.4 ± 0.3 Å at 42, 59, and 79 d, espectively, whereas for He I λ7065 the values are nearly constant r even rising (91 ± 1.4, 106 ± 1.4, and 118 ± 2.2 Å). The same ehaviour as that of He I λ7065 is seen in He I λ5876 but we note hat the Na I D doublet may contaminate the latter line and thus it s not shown here. We interpret the weakening of the He I λ6678 bsorptions as a result of the appearance of H α in emission at a imilar wavelength. Assuming a detached CSM from the progenitor, we conclude that he strong ejecta–CSM interaction initiated sometime prior to + 79 . .2 Flattening in the light cur v es he gri -band light curves are plotted in Fig. 3 , compared with the ype IIb SN 2011dh and the Type Ib iPTF13bvn. The light curves f SN 2019yvr present a characteristic break in the post-maximum ecline rate, leading to a flattening after + 90 d. In Fig. 3 , the break s evident in all optical bands by comparison with SNe that show similar evolution around maximum light. Moreo v er, iPTF13b vn rovided a very good match to the spectroscopic evolution of SN 019yvr around maximum light. Although the spectral match to the N IIb 2011dh is worse, its light curv es pro vide a good match to our bject around maximum light and are more complete than those of PTF13bvn. We interpret the sudden change in the slope of the light curves of N 2019yvr as a result of an extra power source caused by sustained nteraction between the SN ejecta and the CSM. This interpretation s supported by the nearly simultaneous appearance of H α emission n the spectra (see Section 3.1 ). From close inspection of Fig. 3 , we onclude that the flattening in the light curves occurs in all bands etween + 70 and + 90 d with respect to maximum light. This is in ccordance with what was found in Section 3.1 . .3 Properties of the CSM ased on the light curves and spectral ev olution, we ha ve determined hat the interaction power starts dominating the decay power between 70 d to + 90 d (i.e. 75–105 d post our inferred explosion epoch). If e assume the presence of a detached CSM structure, the interaction elay indicates a distance to its inner boundary. By adopting a aximum ejecta velocity of ∼10 000 km s −1 from the bluest extent of he He I λλ6678, 7065 absorption components, this gives a distance of .5–9.1 × 10 15 cm or 0.9–1.3 × 10 5 R . These values are comparable o those obtained for SN 2001em (a SN Ic that showed late phase H α n emission, Chugai & Che v alier 2006 ); and SN 2014C (Milisavljevic t al. 2015 ). If the CSM was expelled by stellar winds with a velocity n the range of 50–100 km s −1 , the mass-loss must have occurred up ntil ∼20–60 yr prior to the explosion (assuming a detached CSM). We further study the properties of the CSM by analysing the H α mission in the nebular phase. We measure the H α luminosity, L H α , o derive the mass-loss rate from (e.g. Kuncarayakti et al. 2018 ) ˙M = 2 L H α H α v wind v 3 shock , here v wind is the wind velocity at which the material was expelled uring the final stages of the star’s evolution, v shock is the velocity of he colliding material, and H α is an efficiency factor that we assume o be 0.01 (Che v alier & Fransson 1994 ). We scaled our + 383 d and + 386 d spectra so that they matched he r -band photometry. They were corrected for extinction and the α fluxes were then measured with splot in IRAF . We discard he measurement from the + 426 d spectrum because it likely uffers from host-galaxy contamination. The resulting fluxes were .77 ± 0.19 × 10 −14 erg s −1 cm −2 from the + 383 d spectrum, and .44 ± 0.44 × 10 −14 erg s −1 cm −2 from the + 386 d spectrum, whose verage yields a flux of 8.10 ± 0.47 × 10 −14 erg s −1 cm −2 . With he distance given in Section 2 , we obtained an H α luminosity of α = 2.1 ± 0.6 × 10 39 erg s −1 . For the shock velocity, we adopted 0 000 km s −1 based on typical values for SNe Ib (Liu et al. 2016 ). ssuming a wind velocity of v wind = 50–100 km s −1 , the derived ass-loss rate range is ∼3 − 7 × 10 −5 M  yr −1 . This is comparable o that of SN 2013df, which also shows a late-phase light curve attening (Maeda et al. 2015 ). If a shock velocity of 2000 km s −1 is onsidered, this translates to an upper limit for the mass-loss rate of 4 − 8 × 10 −3 M  yr −1 . A high density of the CSM may prevent he appearance of [O III ] lines, but would not be enough to produce lectron scattering wings in H α (see Section 4 ). The material could e distributed in a clumpy shell with cloud shocks of about 2000 km −1 and faster, lower density shocks in between, which would be esponsible for the H α wings. MNRASL 529, L33–L40 (2024) L36 L. Ferrari et al. M 4 N a o w S H 2 2 c λ c t a b c t s p c [ i i f s t t b a t c p a f s ( M a f v 2 s C c i 2 e w 5 I i t 6 s e c D [ ( 5 W v t c o t ( T e o e p i ( c ( a a a l p 3 o H r a m h t t c S p M T l H b a t c 5 T D ow nloaded from https://academ ic.oup.com /m nrasl/article/529/1/L33/7503375 by Turku U niversity - IN AC TIVE user on 18 January 2024 N E BU L A R SPECTRA ebular spectra of SN 2019yvr obtained on + 383 d, + 386 d, nd + 426 d are plotted in Fig. 4 , along with a + 371 d spectrum f SN 2014C and a + 290 d spectrum of iPTF13bvn. These objects ere classified as H-poor SNe Ib based on their early spectroscopy. 6 imilar to SN 2014C, SN 2019yvr developed a narrow ∼2000 km s −1 α emission at late times, although their profiles differ. While in SN 019yvr the line is asymmetric and blue shifted by ∼300 km s −1 , SN 014C showed a compound profile, with one broad ∼1200 km s −1 omponent o v erlapped with narrow ∼250 km s −1 H α and [N II ] λ6548, 6583 components (Milisavljevic et al. 2015 ). The narrow omponents may be linked to CSM material undergoing photoioniza- ion caused by X-rays emitted by the interaction, though these could lso result from contamination by an underlying H II region. The road component is associated with the shock or ejected material olliding with the CSM. In the right panel of Fig. 4 , we rebinned he spectrum of SN 2014C to match the resolution of the FORS pectrum of SN 2019yvr. We conclude that those narrow lines, if resent, are not resolved by our observations. Another difference an be appreciated in the narrow emissions associated with H β and O III ] λλ4959, 5007 that are absent in the case of SN 2019yvr. A striking feature is the strong Ca II near-infrared triplet, which s usually weaker than [O I ] λλ6300, 6364 and [Ca II ] λλ7291, 7324 n SESNe (Jerkstrand et al. 2015 ; Dessart et al. 2021 , 2023a ). This eature is detected in the + 282 d spectrum of SN 2014C, but with a ubstantially weaker intensity (Milisavljevic et al. 2015 ) compared o SN 2019yvr. Typical nebular emission lines are present, such as the aforemen- ioned [O I ], [Ca II ], also Na I D, and [Mg I ] λ4571. The oxygen dou- let shows a double-peaked profile, with a ∼1300 km s −1 blueshift nd a full width at half-maximum of ∼2000 and ∼2500 km s −1 for he bluer and redder component respecti vely, dif ferent from the one- omponent profile in SN 2014C. The Na I D emission presents a broad rofile, with ISM absorption on top. The [Ca II ] and [Mg I ] lines show single component profile. Line identifications are shown in Fig. 4 . The H α profile provides insights into the geometry responsible or such emission. If the CSM was distributed in a spherical structure urrounding the SN, the profile should be box-shaped as in SN 1993J Filippenko, Matheson & Barth 1994 ; Patat, Chugai & Mazzali 1995 ; atheson et al. 2000 ) and SN 2013df (Maeda et al. 2015 ). The bsence of such a profile suggests that the emission may not come rom the outer layers of the ejecta interacting with the CSM, as the elocity should be higher ( ∼10000 km s −1 ; see e.g. Dessart et al. 023b ). The ∼2000 km s −1 width the of H α line indicates that the lower, inner part of the ejecta are interacting with a nearby, dense SM. This CSM may take the form of a circumstellar disc which ould be produced in a binary system, although detailed modelling s required to ascertain this possibility. It has been proposed for SN 014C that the H α emission comes from the interaction between the jecta and a CSM with a torus-like structure (Thomas et al. 2022 ), hich could be also the case for SN 2019yvr. PR O G E N I TO R PR OPERTIES n this section, we aim to constrain the progenitor mass, and then link t with what was obtained in Section 3 . For this purpose, we consider hree methods: (i) the early bolometric light curve modelling, (ii) theNRASL 529, L33–L40 (2024) Note that some authors suggested the possible presence of H features in the pectra of iPTF13bvn and SN 2014C (Kuncarayakti et al. 2015 ; Milisavljevic t al. 2015 ). l F W s c omparison of nebular spectra with synthetic nebular spectra from essart et al. ( 2023a ), and (iii) the oxygen mass estimation from the O I ] λλ6300, 6364 flux following the procedure of Jerkstrand et al. 2014 ). Our results are summarized in Fig. 5 . .1 Hydrodynamical model e used the 1D Lagrangian hydrodynamic code (Bersten, Ben- enuto & Hamuy 2011 ) to model the bolometric light curve and he photospheric velocity evolution of SN 2019yvr. As initial onfigurations for our hydrodynamical models, we adopted He stars f different masses from Nomoto & Hashimoto 1988 , which follow he complete evolution of the stars with Zero Age Main Sequence ZAMS) masses of 13, 15, 18, and 25 M , to the pre-SN conditions. he simulations’ free parameters are the explosion energy ( E ), the jected mass ( M ej ), the mass of synthesized 56 Ni ( M Ni ), and the extent f outward mixing of 56 Ni (as a fraction of the pre-SN mass). The nergy is deposited at a certain mass coordinate, M cut , within the re-SN structure. It is assumed that the matter inside M cut collapses nto a compact remnant while the outer mass is ejected. We computed the bolometric light curve for SN 2019yvr based on g –r ) and ( r –i ) colour curves using the bolometric-correction versus olour calibrations for SNe Ib given by Lyman, Bersier & James 2014 ; see their table 2). We applied extinction and distance values s given in Section 2 to derive bolometric luminosities, and then veraged the results obtained from both colour indices. Finally, to pproximate photospheric velocities, we measured the Fe II λ5169 ine velocity from the location of the absorption minimum. Fig. 5 , panels (a) and (b), show the results of the modelling. Our referred model corresponds to a pre-SN model with a mass of .3 M , E = 4 × 10 50 erg, M Ni = 0.088 M  and an e xtensiv e mixing f 0.93. We also assume a M cut = 1.5 M , leading to an M ej = 1.8 M . o we ver a model with a pre-SN mass of 4 M  also produces a easonable match to the data. Therefore we propose progenitors with pre-SN mass between 3 and 4 M , which corresponds to a ZAMS ass of 13–15 M . We have also tested models with higher masses which require igher energy in order to reproduce the e xpansion v elocities, leading o worse fitting to the light curve [see Fig. 5 , panels (a) and (b)]. For hese more massive progenitors, we found that no set of parameters an fit both the bolometric light curve and the velocities together. pecifically, in Fig. 5 we show two models corresponding to 8.0 M  re-SN models with E = 1 foe and E = 5 foe, and equal values for ej = 6.2 M  ( M cut = 1.8 M ), M Ni = 0.1 M , and mixing = 0.98. he first case reproduces well the velocities but not the bolometric ight curve, and vice versa. Since the lowest pre-SN mass model provided by Nomoto & ashimoto ( 1988 ) is that of 3.3 M , we are not able to model the olometric light curve and photospheric velocities for lower masses nd constraint the inferior limit of the progenitor star mass with his approach. The derived ZAMS mass between 13 and 15 M  thus orresponds to an upper limit. .2 Model nebular spectra and the [O I ]/[Ca II ] ratio he flux ratio of nebular [O I ] λλ6300, 6364 to [Ca II ] λλ7291, 7324 ines has been suggested as an indicator of the pre-SN mass (e.g. ransson & Che v alier 1989 ; Maeda et al. 2007 ; Fang et al. 2022 ). e calculated these ratios on the + 383 and the + 426 d EFOSC2 pectra by fitting Gaussian profiles and subtracting the strong local ontinuum. In the case of [O I ], we used two Gaussians to account SN 2019yvr L37 Figure 4. Left panel: nebular spectra of SN 2019yvr compared to the transitional object SN 2014C (Milisavljevic et al. 2015 ), and to its best match in the early phase, iPTF13b vn (K uncarayakti et al. 2015 ). Spectra are corrected by redshift and not corrected by extinction. Prominent emission features are identified and labelled. The telluric A -band in SN 2019yvr is marked with a shadowed region. Right panel: H α and oxygen doublet profiles of SN 2019yvr ( + 386 d, blue) and SN 2014C ( + 371 d, pink). The spectrum of SN 2014C has been rebinned in both panels to match the resolution of the SN 2019yvr spectrum. Both interacting events transition to a SN IIn-like spectrum, showing strong H α emission. H α emission lines in iPTF13bvn are not associated with the SN ejecta but with an underlying H II region. Zero velocities are taken at 6300 and 6563 Å. Figure 5. Progenitor mass estimation summary. Panels (a) and (b) show the bolometric light curve and velocity measurements of SN 2019yvr (stars) together with the output of the hydrodynamical models for stars with final He core masses of 3.3, 4.0, and 8.0 M  (solid, dotted, and dashed lines, respectively; see Section 5.1 ). The preferred model is the least massive. In panel (c), we plot the evolution of oxygen to calcium flux ratio from synthetic spectra (diamonds; Dessart et al. 2023a ) and our measurements for SN 2019yvr (stars), which fall in the region between 3.0 and 3.5 M  He core masses (see Section 5.2 ). Panel (d) shows oxygen yields from Nomoto et al. ( 1997 ), Limongi & Chieffi ( 2003 ), and Rauscher et al. ( 2002 ) for different zero-age main sequence masses and the oxygen core mass derived from [O I ] doublet flux (see Section 5.3 ). f ( w f [ c a S y i n c f D ow nloaded from https://academ ic.oup.com /m nrasl/article/529/1/L33/7503375 by Turku U niversity - IN AC TIVE user on 18 January 2024or the flux, as the line profile is not well fitted by one Gaussian only see Fig. 4 ; right panel). We use the grid of models published in Dessart et al. ( 2023a ), here the spectral evolution between 100 and 400 d is calculated or a wide range of initial He masses. We measured the flux ratio O I ]/[Ca II ] on these models by fitting a single Gaussian profile entred at 6300 Å for [O I ] and at 7304 Å for [Ca II ], since the doubletsre blended. We compared these results to our measurements in N 2019yvr at both epochs, as in Fig. 5 panel (c). This approach ields a progenitor with a helium mass between 3.0–3.5 M , which s in agreement with the result from the hydrodynamical model. We ote, ho we ver, that the line fluxes in the model spectra – which are omputed without CSI – are substantially smaller than those obtained or SN 2019yvr. This difference can be due to a contribution fromMNRASL 529, L33–L40 (2024) L38 L. Ferrari et al. M t t 5 F e fl f s m ( 2 t W a f d h a t e t e t F i 6 W s s fl a i s o e i S a 2 b m e e i K 2 l m S s t s B C 1 s w g s A T r b a D a o J J b n p N w U A O I l C b a i 0 t a s c o N a ( p o f A a u t S 2 M f C P C t A u p 0 M d U D ow nloaded from https://academ ic.oup.com /m nrasl/article/529/1/L33/7503375 by Turku U niversity - IN AC TIVE user on 18 January 2024he CSI in the line fluxes. We cannot ascertain how this may affect he flux ratios. Therefore, this caveat should be kept in mind. .3 [O I ] doublet flux ollowing the procedure described in Jerkstrand et al. ( 2014 ), we stimate the oxygen core minimum mass responsible for [O I ] doublet ux emission. As in Section 5.2 , we assume that the flux comes only rom the ejecta and has no contribution from the CSI. This is a trong assumption, since it has been shown that at late times the aterial excited by CSI can play a major role in the spectral features Dessart et al. 2023b ). Furthermore, the models by Dessart et al. 023a without CSI have much lower lines fluxes, as do the spectra of he SNe that do not appear to have CSI (SN 2011dh and iPFT13bvn). e thus consider this an upper limit for the pre-SN progenitor mass, nd leave further analysis on how the CSI may affect the spectral eatures for the accompanying paper. The flux measurement is performed as detailed in Section 4 in the ereddened spectrum at + 383 d, due to its high quality and minimal ost contamination. Temperature estimation from [O I ] λ5577 is not vailable due to the absence of the line. We therefore assume a typical emperature for these regions of 3000 K. In these conditions, the stimated core oxygen mass is ∼1.1 ± 0.3 M . If we assume a higher emperature of 3500 K, the oxygen mass drops to ∼0.4 ± 0.09 M . In both cases, following oxygen production yields from Nomoto t al. ( 1997 ), Rauscher et al. ( 2002 ), and Limongi & Chieffi ( 2003 ), he estimates indicate a progenitor mass between 15 and 20 M  [see ig. 5 ; panel (d)]. This value is somewhat higher than those obtained n Sections 5.1 and 5.2 . C O N C L U S I O N S e have presented light curves and spectra of SN 2019yvr that how clear signatures of late-time interaction with a CSM. Time- eries observations allowed us to constrain the onset of light curve attening and H α emission line. We estimated the timing of the CSI nd thus the CSM distance to the progenitor, as ∼6.5–9.1 × 10 15 cm n case it is detached from the SN progenitor star. Assuming a teady 50–100 km s −1 wind velocity, this implies a mass-loss rate f ∼3–7 × 10 −5 M  yr −1 , occurring up until ∼20–60 yr prior to the xplosion. Our analysis on the progenitor mass presented in Section 5 is n contradiction with progenitors with pre-SN masses of ≥8 M . uch a star may have started as a single massive star on the ZAMS nd lost the outer layers via vigorous winds, but in the case of SN 019yvr a less-massive star that lost its H-rich envelope through inary interactions (e.g. Fang et al. 2019 ; Drout et al. 2023 ) is a ore plausible scenario. It is also possible that the progenitor star xperienced a hybrid mass-loss mechanism as that discussed by Fang t al. ( 2019 ) and Sun, Maund & Crowther ( 2023 ). The main question lies in how a progenitor with no hydrogen, as ndicated by the early spectra, can lead to a H-rich SN at later times. ilpatrick et al. ( 2021 ) suggested two progenitor scenarios for SN 019yvr: a massive star that went through a series of eruptions in a uminous blue variable (LBV) phase, or a binary system that led to ass-loss episodes timed years to decades ahead of core collapse. un et al. ( 2022 ) suggested a hot and compact progenitor in a binary ystem with a cool and inflated YHG companion. Compared with hese works, our results are compatible with the binary progenitor cenario and do not fa v our a single star going through an LBV phase. y studying the host stellar cluster of SN 2014C, Sun, Maund & rowther ( 2020 ) suggested that the progenitor could have been anNRASL 529, L33–L40 (2024) 1 M  star depleted by binary interaction. Otherwise, a single star hould have retained its hydrogen-rich envelope to be consistent ith the clusters’ inferred age. If this is correct, both SNe must have one through similar evolutionary and mass-loss paths, resulting in tripped progenitors with initial masses well below 20 M . C K N OW L E D G E M E N T S he Finnish National Agency for Education (EDUFI) supported this esearch project through an EDUFI Fellowship. HK was funded y the Research Council of Finland projects 324504, 328898, nd 353019. MDS is funded by the Independent Research Fund enmark (IRFD) via Project 2 grant 10.46540/2032-00022B. KM cknowledges support from the Japan Society for the Promotion f Science (JSPS) KAKENHI grant (JP20H00174) and by the SPS Open Partnership Bilateral Joint Research Project between apan and Finland (JPJSBP120229923). NUTS is funded in part y the Instrument center for Danish Astrophysics (IDA). LCOGT etwork data were obtained through OPTICON (PI Stritzinger, rogramme ID 2020A/031) and NOAO (PI Bose, programme ID OAO2020A-017) time allocations. Based on observations made ith the Nordic Optical Telescope, owned in collaboration by the niversity of Turku and Aarhus University, and operated jointly by arhus University, the University of Turku and the University of slo, representing Denmark, Finland and Norway, the University of celand and Stockholm University at the Observatorio del Roque de os Muchachos, La Palma, Spain, of the Instituto de Astrofisica de anarias. Data were also obtained with ALFOSC, which is provided y the Instituto de Astrofisica de Andalucia (IAA) under a joint greement with the University of Copenhagen and NOT. This work ncludes data collected at ESO via programme IDs 1103.D-0328, 105.D-0511, and 1106.D-0811. We thank the Subaru staff for he data taken by the Subaru Telescope (S19B-054); the authors cknowledge the very significant cultural role and reverence that the ummit of Maunakea has always had within the indigenous Hawaiian ommunity. We are most fortunate to have the opportunity to conduct bservations from this mountain. YZC is supported by the National atural Science Foundation of China (NSFC,; grant no. 12303054) nd the International Centre of Supernov ae, Yunnan K ey Laboratory no. 202302AN360001). AP is supported by the PRIN-INAF 2022 roject ‘Shedding light on the nature of gap transients: from the bservations to the models’. LG acknowledges financial support rom the Spanish Ministerio de Ciencia e Innovaci ´on (MCIN), the gencia Estatal de Investigaci ´on (AEI) 10.13039/501100011033, nd the European Social Fund (ESF) ‘Investing in your future’ nder the 2019 Ram ´on y Cajal programme RYC2019-027683-I and he PID2020-115253GA-I00 HOSTFLOWS project, from Centro uperior de Investigaciones Cient ´ıficas (CSIC) under the PIE project 0215AT016, and the programme Unidad de Excelencia Mar ´ıa de aeztu CEX2020-001058-M. CPG acknowledges financial support rom the Secretary of Universities and Research (Go v ernment of atalonia) and by the Horizon 2020 Research and Innovation rogramme of the European Union under the Marie Skłodowska- urie and the Beatriu de Pin ´os 2021 BP 00168 programme, from he Spanish Ministerio de Ciencia e Innovaci ´on (MCIN) and the gencia Estatal de Investigaci ´on (AEI) 10.13039/501100011033 nder the PID2020-115253GA-I00 HOSTFLOWS project, and the rogramme Unidad de Excelencia Mar ´ıa de Maeztu CEX2020- 01058-M. TEMB acknowledges financial support from the Spanish inisterio de Ciencia e Innovaci ´on (MCIN), the Agencia Estatal e Investigaci ´on (AEI) 10.13039/501100011033, and the European nion Next Generation EU/PRTR funds under the 2021 Juan de la SN 2019yvr L39 C I C p 0 F C a U s I D N 2 R A B B C C C D D D D D D E F F F F F F J J K K K K L L L M M M M M M M M N N N P P R R S S S S S S S S S S S S S T T T T V W Y Y S S s P o A c 1 A 2 L 3 U 4 F 5 2 6 1 7 S 8 ( 9 S 1 b D ow nloaded from https://academ ic.oup.com /m nrasl/article/529/1/L33/7503375 by Turku U niversity - IN AC TIVE user on 18 January 2024ierva programme FJC2021-047124-I and the PID2020-115253GA- 00 HOSTFLOWS project, from Centro Superior de Investigaciones ient ´ıficas (CSIC) under the PIE project 20215AT016, and the rogramme Unidad de Excelencia Mar ´ıa de Maeztu CEX2020- 01058-M. SM acknowledges support from the Research Council of inland project 350458. 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Instituto de Astrof ´ısica de La Plata, CONICET, B1900FWA, La Plata, rgentina Facultad de Ciencias Astron ´omicas y Geof ´ısicas Universidad Nacional de a Plata, Paseo del Bosque S/N B1900FWA, La Plata, Argentina Kavli Institute for the Physics and Mathematics of the Universe (WPI), The niversity of Tokyo, Kashiwa, 277-8583 Chiba, Japan Department of Physics and Astronomy, University of Turku, FI-20014 Turku, inland Finnish Centre for Astronomy with ESO (FINCA), University of Turku, FI- 0014, Finland Department of Physics and Astronomy, Aarhus University, Ny Munkegade 20, DK-8000 Aarhus C, Denmark Department of Astronomy, Kyoto Univer sity, Kitashirakawa-Oiwak e-cho, akyo-ku, Kyoto 606-8502, Japan Inter disciplinary Theor etical and Mathematical Sciences Program iTHEMS), RIKEN, Wako, Saitama 351-0198, Japan Department of Physics, University of California, Berkeley, CA 94720, United tates of America 0 Institut d’Astrophysique de Paris, CNRS-Sorbonne Universit ´e, 98 bis oule vard Ar ago, F-75014 Paris, Fr ance MNRASL 529, L33–L40 (2024) L40 L. Ferrari et al. MNRASL 529, L33–L40 (2024) 11 The Oskar Klein Centre, Department of Astronomy, Stockholm University, AlbaNova, SE-10691 Stockholm, Sweden 12 Max-Planck-Institut f ¨ur Astrophysik, Karl-Sc hwarzsc hild-Str. 1, D-85748 Garching, Germany 13 Astrophysics Research Institute, Liverpool John Moores University, 146 Brownlow Hill, Liverpool L3 5RF, UK 14 Mets ¨ahovi Radio Observatory, Aalto University , Mets ¨ahovintie 114, FI- 02540 Kylm ¨al ¨a, Finland 15 Department of Electronics and Nanoengineering , Aalto Univer sity , P.O. BOX 15500, FI-00076 AALTO, Finland 16 Department of Physics, Virginia Tech, Blac ksb urg, VA 24061, USA 17 Department of Astronomy, The Ohio State University, 140 W. 18th Avenue, Columbus, OH 43210, USA 18 Yunnan Observatories, Chinese Academy of Sciences, Kunming 650216, P.R. China 19 Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650216, P.R. China 20 International Centre of Superno vae , Yunnan Key Laboratory, Kunming 650216, P.R. China 21 INAF - Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy 22 INAF - Osservatorio Astronomico di Brera, Via E. Bianchi 46, I-23807 Merate (LC), Italy 23 European Southern Observatory, Alonso de C ´ordova 3107, Casilla 19, Santiago, Chile 24 Graduate Institute of Astronomy, National Central University, 300 Jhongda Road, 32001 Jhongli, Taiwan 25 Institute of Space Sciences (ICE-CSIC), Campus UAB, Carrer de Can Magrans, s/n, E-08193 Barcelona, Spain 26 Institut d’Estudis Espacials de Catalunya (IEEC), E-08034 Barcelona, Spain 27 Astronomical Observatory, University of Warsaw, Al. Ujazdowskie 4, PL- 00-478 Warszawa, Poland 28 School of Physics and Astronomy, Cardiff University, Queens Buildings, The Parade, Cardiff CF24 3AA, UK 29 School of Sciences, European Univer sity Cyprus, Dio g enes Street, Engomi, 1516, Nicosia, Cyprus 30 Astrophysics Resear ch Centr e, School of Mathematics and Physics, Queens University Belfast, Belfast BT7 1NN, UK 31 Instituto de Alta Investigaci ´on, Universidad de Tarapac ´a, Casilla 7D, 1001236, Arica, Chile 32 Astrophysics Resear ch Centr e, School of Mathematics and Physics, Queen’s University Belfast, Belfast BT7 1NN, UK This paper has been typeset from a T E X/L A T E X file prepared by the author. © The Author(s) 2023. Published by Oxford University Press on behalf of Royal Astronomical Society. 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