Flash ionization signatures in the type Ibn supernova SN 2019uo
Anjasha Gangopadhyay, Kuntal Misra, Daichi Hiramatsu, Shan-Qin Wang, Griffin Hosseinzadeh, Xiaofeng Wang, Stefano Valenti, Jujia Zhang, D. Andrew Howell, Iair Arcavi, G.C. Anupama, Jamison Burke, Raya Dastidar, Koichi Itagaki, Brajesh Kumar, Brijesh Kumar, Long Li, Curtis McCully, Jun Mo, Shashi Bhushan Pandey, Craig Pellegrino, Hanna Sai, D.K. Sahu, Pankaj Sanwal, Avinash Singh, Mridweeka Singh, Jicheng Zhang, Tianmeng Zhang, Xinhan Zhang
DDraft version December 18, 2019
Typeset using L A TEX twocolumn style in AASTeX63
Flash ionization signatures in the type Ibn supernova SN 2019uo
Anjasha Gangopadhyay,
1, 2
Kuntal Misra,
1, 3
Daichi Hiramatsu,
4, 5
Shan-Qin Wang, Griffin Hosseinzadeh, Xiaofeng Wang, Stefano Valenti, Jujia Zhang, D. Andrew Howell,
4, 5
Iair Arcavi,
10, 11
G.C. Anupama, Jamison Burke,
4, 5
Raya Dastidar,
1, 13
Koichi Itagaki, Brajesh Kumar, Brijesh Kumar, Long Li, Curtis McCully,
4, 5
Jun Mo, Shashi Bhushan Pandey, Craig Pellegrino,
4, 5
Hanna Sai, D.K. Sahu, Pankaj Sanwal,
1, 2
Avinash Singh,
12, 15
Mridweeka Singh,
16, 1
Jicheng Zhang, Tianmeng Zhang, andXinhan Zhang Aryabhatta Research Institute of observational sciencES, Manora Peak, Nainital 263 002 India School of Studies in Physics and Astrophysics, Pandit Ravishankar Shukla University, Chattisgarh 492 010, India Department of Physics, University of California, 1 Shields Ave, Davis, CA 95616-5270, USA Las Cumbres Observatory, 6740 Cortona Drive Suite 102, Goleta, CA, 93117-5575 USA Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA Guangxi Key Laboratory for Relativistic Astrophysics, School of Physical Science and Technology, Guangxi University, Nanning 530004,People’s Republic of China Center for Astrophysics | Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138-1516, USA Physics Department and Tsinghua Center for Astrophysics, Tsinghua University, China Yunnan Astronomical Observatory of China, Chinese Academy of Sciences, Kunming, 650011, China The School of Physics and Astronomy, Tel Aviv University, Tel Aviv 69978, Israel CIFAR Azrieli Global Scholars program, CIFAR, Toronto, Canada Indian Institute of Astrophysics, Koramangala 2nd Block, Bangalore 560034, India Department of Physics & Astrophysics, University of Delhi, Delhi-110 007 Itagaki Astronomical Observatory, Japan Joint Astronomy Programme, Department of Physics, Indian Institute of Science, Bengaluru 560012, India Korea Astronomy and Space Science Institute, 776 Daedeokdae-ro, Yuseong-gu, Daejeon 34055, Republic of Korea National Astronomical Observatory of China, Chinese Academy of Sciences, Beijing, 100012, China
ABSTRACTWe present photometric and spectroscopic observations of the type Ibn supernova (SN) 2019uo, thesecond ever SN Ibn with flash ionization (He II , C III , N
III ) features in its early spectra. SN 2019uodisplays a rapid post-peak luminosity decline of 0.1 mag d − similar to most of the SNe Ibn, but isfainter ( M Vmax = − . ± .
24 mag) than a typical SN Ibn and shows a color evolution that placesit between SNe Ib and the most extreme SNe Ibn. SN 2019uo shows P-cygni He I features in theearly spectra which gradually evolves and becomes emission dominated post peak. It also shows fasterevolution in line velocities as compared to most other members of the type Ibn subclass. The bolometriclight curve is fairly described by a Ni + circumstellar interaction model.
Keywords: supernovae: general – supernovae: individual: SN 2019uo – galaxies: individual: – tech-niques: photometric – techniques: spectroscopic INTRODUCTIONSupernovae (SNe) undergoing interaction with a cir-cumstellar medium (CSM) provide a unique window inthe evolutionary phases of stars. Interaction, in general,produces narrow emission lines — broader than H II Corresponding author: Anjasha [email protected], [email protected] regions but narrower than lines arising from the outerejecta of the supernova (Hosseinzadeh et al. 2017). How-ever, in some cases interaction happens below the pho-tosphere without any observable narrow emission lines(e.g. Morozova et al. 2017; Andrews & Smith 2018).SNe IIn (Schlegel 1990) and SNe Ia-CSM display nar-row H lines indicative of interaction with a H-rich CSM.Approximately 1% of core-collapse SNe (CCSNe) showlittle H and narrow He features ( ∼ − ). Withthe discovery of SN 2006jc, Pastorello et al. (2007) in- a r X i v : . [ a s t r o - ph . H E ] D ec Gangopadhyay et al. troduced this class as SNe Ibn, whose spectral featuresshow interaction signatures between SN ejecta and aHe-rich CSM. This is defined in analogy with SNe IIn,which show narrow H features (Schlegel 1990). SNe thatare embedded in dense CSM may also show short-livednarrow high ionization emission lines ( ≤
10 days) owingto the recombination of the CSM following the shockbreakout flash. These features are known as “flash fea-tures” (e.g. Gal-Yam et al. 2014). Hosseinzadeh et al.(2017) analysed a sample of SN Ibn light curves andshowed that unlike SNe IIn, SNe Ibn are rather uni-form in their light curve shape with rapid decay rates of0.05–0.15 mag d − . SNe Ibn may have double-peakedlight curves like SNe IIn, but they show a faster risethan SNe IIn (Hosseinzadeh et al. 2017). On the otherhand, Pastorello et al. (2016) showed that the classis heterogeneous with many outliers: OGLE-2012-SN-006 (Pastorello et al. 2015a) has a very slow decline;LSQ13ccw (Pastorello et al. 2015b) is faint and fast-declining; SNe 2005la and 2011hw (Pastorello et al.2015c) are transitional type IIn/Ibn events; SN 2010al(Pastorello et al. 2015c) is the earliest detected SN Ibnwith a slow rise and decline. Karamehmetoglu et al.(2019) recently identified a rapid evolving SN 2018bcc.SNe Ibn have bluer continuum than other CCSNe. SomeSNe Ibn show P Cygni He I emission, while others tran-sition from narrow to intermediate-width He I emissions(Hosseinzadeh et al. 2017).So far, only indirect progenitor constraints for SNe Ibnare available. Pastorello et al. (2007) suggest Wolf-Rayet(WR) H-free atmospheres generate the He-rich CSM.The best studied case for unstable mass loss from a WRprogenitor is SN 2006jc, for which an optical transientwas detected at the SN location two years prior to explo-sion (Foley et al. 2007; Pastorello et al. 2007; Smith et al.2008). Alternatively, CSM can be produced by strip-ping material from envelopes of massive binaries (Foleyet al. 2007). However, a low-mass progenitor has beensuggested for PS1-12sk, which occurred in a non-star-forming host (Sanders et al. 2013; Hosseinzadeh et al.2019) — unlikely for a CCSN ( ≤ HST images.In this paper we study the evolution of one such typeIbn SN 2019uo which was discovered on 2019 January17.8 UT (JD 2458501.3) by Koichi Itagaki at R . A . =12 h m . s , Decl . = +41 ◦ (cid:48) (cid:48)(cid:48) (J2000.0). The SNlocation is 0 . (cid:48)(cid:48) . (cid:48)(cid:48) V , He II and O V . However, this classification oftype II SN was modified later by Fremling et al. (2019)and SN 2019uo was classified as a type Ibn. Prominentnarrow emission lines of He I in the initial spectra ofSN 2019uo indicating a P-cygni velocity of 650 km s − justified the type Ibn classification. SN 2019uo is thesecond SN Ibn to show these features after SN 2010al.Adopting H = 73 km s − Mpc − , we obtain a lumi-nosity distance of 88.8 Mpc for SN 2019uo. The MilkyWay extinction along the line of sight of SN 2019uo is A V = 0 .
035 mag (Schlafly & Finkbeiner 2011). For es-timating the extinction along the line of sight withinhost galaxy, we estimate equivalent widths of the Na I Dline in the first three spectra of SN 2019uo. Using theformulation by Munari & Zwitter (1997) and Poznanskiet al. (2012), we estimate A V = 0 . B − V colors of SN 2019uo intoclose agreement with SNe 2006jc and 2010al. Thus, weadopt a total A V = 0 .
287 mag. The temporal and spec-tral evolution of SN 2019uo and the detailed modelingof the bolometric light curve is discussed in the sectionsto follow. DATA ACQUISITION AND REDUCTIONWe observed SN 2019uo with Las Cumbres Obser-vatory (LCO) in the
UBVgri filters from ∼ UBVRI/ugri were alsotaken with 0.7m BITRAN-CCD Imaging System locatedin Japan; 0.8m Tsinghua-NAOC Telescope (TNT),Xinglong Observatory, China; 1.04m SampurnanandTelescope (ST); 1.30m Devasthal Fast Optical Tele-scope (DFOT), ARIES, India; 2.00m Himalayan Chan-dra Telescope (HCT), IAO, Hanle, India and Lijiang2.4m Telescope (LJT), Yunnan Observatories (YNAO),China. We performed image subtraction using HighOrder Transform of PSF ANd Template Subtraction(HOTPANTS) (Becker 2015). The instrumental mag-nitudes were estimated using IRAF (Tody 1986, 1993)and DAOPHOT (Stetson 1987). The LCO photometrywas done using lcogtsnpipe (see Valenti et al. 2011, https://github.com/acbecker/hotpants Image Reduction and Analysis Facility Dominion Astrophysical Observatory Photometry https://github.com/svalenti/lcogtsnpipe lash ionization signatures in type Ibn supernova SN 2019uo Table 1.
Photometry of SN 2019uo
Date JD Phase † U B g V r i
Telescope(yyyy-mm-dd) (2458000+) (day) (mag) (mag) (mag) (mag) (mag) (mag)2019-01-18 501.8 -6.8 — 17.865 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± † with respect to JD max = 2458508 . gri bands and the Landolt standard fields taken onthe same night by the same instrument as the scienceimages for UBV . Wherever required, the RI magnitudeswere converted to ri using the equations of Jordi et al.(2006). The photometry of SN 2019uo is presented inTable 1.The spectroscopic observations were taken at 9 epochsspanning up to ∼
88 days after discovery. The 1Dwavelength- and flux-calibrated spectra were extractedusing the floydsspec pipeline (Valenti et al. 2014) for theLCO data. Spectroscopic data reduction of the 2.2 mand 2.4 m telescopes was done using the APALL task inIRAF followed by wavelength and flux calibration. Theslit loss corrections were done by scaling the spectra tothe photometry. Finally, the spectra were corrected forthe heliocentric redshift of the host galaxy. The log ofspectroscopic observations is given in Table 2. PHOTOMETRIC EVOLUTION OF SN 2019UOThe complete multi-band light curve of SN 2019uois shown in Figure 1. With our available observations,we were able to trace the epoch of maximum in all the bands. The date of maximum and its brightnesswere determined by fitting a cubic spline to the
UB-Vgri light curves. The maximum in r -band occurredon JD 2458508 . ± . . ± .
03 mag. The errors reported are obtainedfrom interpolated measurements around the peak. Weuse days since r -maximum ( r max ) as a reference epochthroughout the paper. Assuming that the discoverydate is close to explosion, we estimated a rise timeof 8 . ± . r -band light curve, between 0–36 days, decayswith a rate of 0 . ± .
005 mag d − . The g , B , V and i bands follow approximately the same decline rate.The sample of SNe Ibn in Hosseinzadeh et al. (2017) arefast-evolving with a typical decline rate of 0.1 mag d − during the first month post-maximum. SN 2019uo fol-lows the same decline rate.Figure 2 shows the absolute magnitude light curve ofSN 2019uo along with other SNe Ibn after correcting fordistance and extinction. The peak r -band absolute mag-nitude of SN 2019uo is − . ± .
24 mag, which is at thefainter end of SN Ibn sample. The blue band in Figure 2shows the average light curve (comprising of 95% of theSN Ibn data) of SNe Ibn taken from Hosseinzadeh et al.(2017). The average light curve was generated by using
Gangopadhyay et al.
Table 2.
Log of spectroscopic observations of SN 2019uo.
Date JD − † Telescope Instrument Range (˚A)2019-01-19 503.4 -5.2 2.4 m LJT YFOSC 3500-88002019-01-20 503.9 -4.7 2.0 m FTN FLOYDS 3200-90002019-01-21 504.9 -3.7 2.0 m FTN FLOYDS 3200-90002019-01-23 506.9 -1.7 2.0 m FTN FLOYDS 3200-90002019-01-24 508.3 -0.3 2.4 m LJT YFOSC 3500-88002019-01-28 512.4 3.8 2.0 m FTN FLOYDS 3200-90002019-02-05 519.9 11.3 2.0 m FTN FLOYDS 3200-90002019-02-08 523.8 15.2 2.2 m China BFOSC 4000-100002019-02-14 529.2 20.6 2.4 m LJT YFOSC 3500-8800 † with respect to JD max = 2458508 .
10 0 10 20 30 40
Days Since r max A pp a r e n t M a g n i t u d e i-2.0r-1.0Vg+1.0B+1.5U+3.0 Figure 1.
UBVgri light curve evolution of SN 2019uo. a Gaussian process to fit a smooth curve to the com-bined light curves on the sample of Hosseinzadeh et al.(2017). The fit was performed in log-log space to en-sure consistency and smoothness between the early andlate time light curves. The average light curve, thus,generated also uses the Gaussian process to fit positiveand negative residuals. It is to note that SN 2019uois ∼ B − R/r color evolution of SN 2019uowith a number of type Ibn SNe, which usually showheterogeneity in their color evolution. The B − r colorof SN 2019uo increases up to 0.64 mag ∼
20 days post r max , subsequently becoming blue at ∼
36 days. Sim-ilarly, for SN 2010al and iPTF14aki the B − r colorincreases up to ∼ ∼
30 days post R max . Thus,SN 2019uo shows a color evolution similar to SN 2010aland iPTF14aki. At similar epochs, the color evolutionof SN 2006jc was extremely blue ( − . SPECTRAL EVOLUTIONThe spectral evolution of SN 2019uo from − . lash ionization signatures in type Ibn supernova SN 2019uo A b s m a g ( R / r - B a n d ) Normalised Ibn2006jc2010al 2011hwOGLE12-006iPTF14aki 2015G2015U2019uo max ( B - R / r ) o Figure 2.
R/r -band absolute magnitude light curve and B − R/r colour curve of SN 2019uo. The comparison sampleincludes SNe 2006jc (Pastorello et al. 2007; Foley et al. 2007),2010al (Pastorello et al. 2015b), OGLE-SN-006 (Pastorelloet al. 2015a), 2011hw (Pastorello et al. 2015b), iPTF14aki(Hosseinzadeh et al. 2017), 2015U (Shivvers et al. 2016; Hos-seinzadeh et al. 2017) and 2015G (Hosseinzadeh et al. 2017).
The early spectral sequence shows a unique blue con-tinuum similar to SN 2010al. Blackbody fits to the firstthree spectra ( − . − .
7, and − . ≤
137 km s − ;unresolved) in the early spectrum of SN 2019uo is mostlikely due to interstellar gas in the host galaxy. Promi-nent emission features in the first three spectra ( − . − . ∼ II at 4686 ˚A, whereas the blue com-ponent arises from a blend of C III
III
III feature at 5696 ˚A.Pastorello et al. (2015c) interpreted these as flash ion-ization signatures in a He-rich CSM (also see Gal-Yam
Rest Wavelength ( ) l o g F + c o n s t a n t − − − − − H e O II H e C III / N III H e II H e H e II C III H e H e H e H e H α H e II Figure 3.
Spectral evolution of SN 2019uo from − . max . Prominent He features are seen in theearly spectra. Flash ionization signatures of He II , C III andN
III are also seen.
III features were found in PTF12ldyand iPTF15ul (Hosseinzadeh et al. 2017), SN 2010al isthe only previous SN Ibn where flash ionization signa-tures of C
III and He II , typical of SNe II, are both seen.Cooke et al. (2010) and Silverman et al. (2010) identifiedsuch lines to be originating from a WR wind, previouslynoted in SNe IIn (e.g., SN 1998S; Fassia et al. (2001)and SN 2008fq; Taddia et al. (2013)). We also identifya He II − at − . − . ∼ II andHe II features, respectively.Figure 4 shows the spectra of SNe 1998S (type IIn)and 2010al (type Ibn) in comparison with SN 2019uo.These two SNe have previously shown flash ionizationsignatures. While the spectrum of SN 2010al shows C III features around 4650 ˚A only, SN 2019uo shows C
III
Gangopadhyay et al. l o g F + c o n s t a n t − − − H e H e C III / N III H e II H e H e II C III H e H e H e H e H α H e II Figure 4.
Spectral comparison of SN 2019uo at − . features around 4650 ˚A and at 5696 ˚A. The inset inFigure 4 highlights these features.As the SN evolves further (3.8 days), the narrow He I P Cygni feature is superimposed on a broader base (thecontinuum is not flat). The flash ionization spectral fea-tures vanish completely during this epoch. From 11–21 days, features of Ca II , Si II , and Na I D also startdeveloping (see Figure 3). Figure 5 shows the compari-son of SN 2019uo with a group of SNe Ibn between 3 – 10days after peak. The He I I P Cygni feature of SN 2019uo is narrower, and issuperimposed over a broader emission line. On the otherhand, the He I P-Cygni profile in SN 2010al is over a flatcontinuum. Flash ionization signatures in SN 2010al arestill visible at this phase, but these features have van-ished in SN 2019uo. The line evolution of SN 2019uoshows that it belongs to the “P Cygni” subclass (follow-
Rest Wavelength ( ) l o g F + c o n s t a n t H e H e C III / N III H e II H e H e II C III H e H e H e H e Figure 5.
Comparison of the spectrum of SN 2019uo toother SNe Ibn. SN 2019uo and SN 2010al show distinctnarrow P-cygni He I spectroscopic features. The data forthis are taken from — SNe 2010al (Pastorello et al. 2015c),2011hw (Pastorello et al. 2015c), PTF11rfh (Hosseinzadehet al. 2017), PTF12ldy (Hosseinzadeh et al. 2017), LSQ13ccw(Pastorello et al. 2015b) and iPTF14aki (Hosseinzadeh et al.2017) ing the interpretation of Hosseinzadeh et al. 2017). TheP Cygni He I features are narrow but gradually broadenwith time. The physical explanation behind the originof the “P Cygni” subclass could be a shell of He aroundthe progenitor star surrounded by a dense CSM. As theoptically thick shell is lit by the explosion, the narrowP Cygni features transition to broader emission as theshell is swept up by the SN ejecta. The viewing angledependence could also affect this scenario; if the CSM isasymmetric and we have a He rich torus, then P Cygnifeatures would only be visible if the system is viewededge-on, while emission features can be seen only if it isviewed face-on. However, this scenario was questionedby Karamehmetoglu et al. (2019) which suggested thatHe I line fluxes are largely dependent on density, temper- lash ionization signatures in type Ibn supernova SN 2019uo He 5876 V e l o c i t y ( k m s e c − ) He 6678 He 7065 E W () Days Since R/r max
Figure 6.
Evolution of line velocities and equivalent widths of He I emission lines is shown in top and bottom panels, respectively.The data for this are taken from — SNe 2006jc (Foley et al. 2007; Pastorello et al. 2008), 2010al (Pastorello et al. 2015c), 2011hw(Pastorello et al. 2015c), PTF11rfh (Hosseinzadeh et al. 2017), LSQ12btw (Pastorello et al. 2015b), OGLE12-006 (Pastorelloet al. 2015a), PTF12ldy (Hosseinzadeh et al. 2017), iPTF13beo (Gorbikov et al. 2014), LSQ13ccw (Pastorello et al. 2015b),iPTF14aki (Hosseinzadeh et al. 2017), 2014av (Pastorello et al. 2016), 2014bk (Pastorello et al. 2016), iPTF15akq (Hosseinzadehet al. 2017), ASASSN-15ed (Pastorello et al. 2015d), and 2015G (Hosseinzadeh et al. 2017). ature and optical depths. Karamehmetoglu et al. (2019)suggest that dominance of emission at late phases is notbecause of being optically thin, but because they lackother lines to branch into it. He ionisation and recom-bination are mostly caused by UV and X-ray, occurringat shock boundary, deep in interacting regions. Eventhough most of the emission and the electron scatteringare produced by the ionised region outside the shock, P-cygni features usually originate from optical depths ≤
1. X-rays penetrating further into the P-cygni produc-ing regions will fill in the absorption and lead to emissionfeatures. Thus, this provides an alternative scenario to the transitioning of P-cygni to emission features of He I lines for type Ibn SNe.The measured the expansion velocities and equivalentwidths (EWs) of three neutral He lines (5876, 6678, and7065 ˚A), wherever visible. We fit the emission lines ofHe I using a Gaussian on a linear continuum. The EWis estimated through the integral of the flux normalizedto the local continuum. We do not measure the EW ofthe P-cygni lines. The velocities reported are estimatedfrom the absorption minima of P-cygni profiles. Figure 6shows the evolution of velocity and EW for a sampleof SNe Ibn taken from Hosseinzadeh et al. (2017) withtime. We see that both the line velocities and EW of the Gangopadhyay et al. L u m i n o s i t y ( x ) ( e r g / s e c ) He II (4686 )
Figure 7. He II luminosity of a sample of SNe II, IIn, andIbn (Khazov et al. 2016) with flash ionization signatures.Blue symbols: type IIn, Black symbols: type II (IIb,IIP andIIL), Red symbols: type Ibn. He lines gradually increase with time and the velocityestimates of SN 2019uo lie in the lower range of SNe Ibn.However, SN 2019uo shows a faster evolution in linevelocities, reaching broader emission profiles as seen inthe P-cygni subclass (Hosseinzadeh et al. 2017) while theemission subclass shows very little velocity evolution.To ascertain the origin of the SNe Ibn, we collected asample of 12 SNe II (including SNe IIb and IIP, IIn) andIbn from Khazov et al. (2016) that showed signatures offlash ionization within 10 days of explosion. Since the Hlines are usually contaminated by the host galaxy, we se-lected the relatively unblended He II II lines are much narrower than lines from theSN ejecta, they can serve as a good tool for probing theflash-ionized CSM. When measuring the luminosities,we removed the continuum by fitting a linear function.Figure 7 shows that the typical luminosity of the He II line for SN 2019uo is similar to the type IIn SNe 1998Sand PTF13ast. MODELING THE BOLOMETRIC LIGHT CURVEOF SN 2019UOTo construct the bolometric light curve of SN 2019uo,the measured flux values were corrected for distance andreddening as given in Section 1. Spectral energy dis-tributions (SEDs) were constructed accounting for the (cid:5) (cid:6) (cid:9) (cid:3) (cid:1) (cid:2) (cid:4) (cid:8) (cid:7) (cid:9) (cid:5) (cid:6) (cid:7) (cid:14) (cid:24) (cid:16) (cid:18)(cid:11)(cid:12)(cid:15) (cid:3) (cid:2) (cid:1) (cid:5) (cid:3) (cid:5) (cid:4) (cid:5) (cid:1) (cid:1) (cid:5) (cid:1) (cid:3) (cid:4) (cid:8) (cid:9) (cid:5)(cid:6)(cid:7) (cid:2) (cid:9) (cid:2) (cid:5) (cid:4) (cid:5) (cid:1) (cid:1) (cid:5) (cid:3) (cid:7) (cid:8) (cid:6) (cid:9) (cid:1) (cid:4) (cid:5) (cid:2) (cid:6) (cid:3) (cid:2) (cid:1) (cid:8) (cid:4) (cid:2) (cid:1) (cid:3) (cid:2) (cid:1) (cid:3)(cid:2)(cid:9) (cid:4)(cid:2)(cid:9) (cid:4) (cid:3) (cid:24) (cid:6)(cid:2)(cid:9) (cid:7) (cid:2) (cid:9) (cid:8) (cid:14) (cid:16) (cid:12) (cid:24) (cid:20) (cid:14) (cid:17) (cid:10) (cid:12) (cid:24) (cid:21) (cid:13)(cid:12) (cid:24) (cid:19)(cid:12) (cid:9) (cid:78) (cid:24)(cid:1) (cid:11) (cid:9)(cid:23) (cid:20) (cid:2)(cid:24)
SN 2019uo
Figure 8.
Best-fit light curves of SN 2019uo using a Nimodel. flux coverage between UV to IR bands using the
Su-perBol (Nicholl 2018) code. The lack of UV and NIRdata was supplemented by extrapolating the SEDs us-ing the blackbody approximation and direct integrationmethod as described in Lusk & Baron (2017). A linearextrapolation was performed in UV regime at late times.The estimated peak bolometric luminosity of SN 2019uois 8 . × erg s − . We used different models to fitthe bolometric light curve at a fixed optical opacity of0.1 cm g − . A Markov Chain Monte Carlo (MCMC)technique was used to obtain the best-fit parameters. Ni model:
Assuming that the peak bolometric lu-minosity is powered by the decay of Ni to Co, wefit the bolometric light curve using Ni model (Arnett1982, 1980). The parameters of the Ni model are theejecta mass M ej , the initial scale velocity of the ejecta v sc0 , the Ni mass M Ni , the gamma-ray opacity of Nidecay photons κ γ, Ni and explosion time t expl . The ini-tial kinetic energy of the ejecta is neutrino-driven and isconsidered to be E k = 0 . M ej v . The best-fit param-eters are tabulated in Table 3 and the best-fit model isdisplayed Fig. 8. The corner plot showing the covarianceof the estimated parameters are represented in Fig. 9.We note that the Ni mass obtained from the poweringmechanism of Arnett (1982) are in concordance with thevalues quoted for several stripped envelope SNe (Lymanet al. 2016; Prentice et al. 2016, 2019). Although the Ni mass inferred from the model is ∼ (cid:12) whichis comparable to that of normal CCSNe, the opacity forthe gamma ray κ γ, Ni emitted from the cascade decay of Ni is 0.01 cm g − , which is significantly smaller thanthe canomical lower limit which is 0.025-0.027 cm g − . lash ionization signatures in type Ibn supernova SN 2019uo Figure 9.
The corner plot of the Ni model displaying covariance of estimated parameters.
Therefore, the Ni model is not a good model in ex-plaining the light curve of SN 2019uo and other modelsmust be employed.
The CSI model and the Ni + CSI model:
Thenarrow He emission lines appearing in the spectra ofSN 2019uo indicate a potential source of circumstellarinteraction (CSI) with a nearby He-rich shell. Thus, thenearby He-rich wind or shell surrounding the progeni- tor could be the essential powering source of the bolo-metric light curve of SN 2019uo. We take into accountthe ejecta-CSM interaction model (i.e., the CSI model)(Chevalier 1982; Chevalier & Fransson 1994; Chugai &Danziger 1994; Ginzburg & Balberg 2012; Liu et al.2018) and the Ni + CSI model (Chatzopoulos et al.2012). To fit the bolometric light curve of SN 2019uo,we adopt the formulation given in Wang & Li (2019).0
Gangopadhyay et al. (cid:6) (cid:1) (cid:3)(cid:2) (cid:1) (cid:5) (cid:4)(cid:5) (cid:1) (cid:2)(cid:7) (cid:1)(cid:3) (cid:4)(cid:8)(cid:9)(cid:5)(cid:6)(cid:7)(cid:2)(cid:9) (cid:3)(cid:5) (cid:1) (cid:2)(cid:7) (cid:3)(cid:7)(cid:8)(cid:6)(cid:9)(cid:1)(cid:4)(cid:5)(cid:2)(cid:6)(cid:9) (cid:4)(cid:7)(cid:5)(cid:7)(cid:5)(cid:7)(cid:6)(cid:1)(cid:1)(cid:2) (cid:2) (cid:2) (cid:3) (cid:3)(cid:3) (cid:3) (cid:3)(cid:3) (cid:3)(cid:3) (cid:3)(cid:3) (cid:3)(cid:3)(cid:3) (cid:3)(cid:3) (cid:3)(cid:3) (cid:3) (cid:7) (cid:11) (cid:9) (cid:34) (cid:22) (cid:24) (cid:15) (cid:16) (cid:20) (cid:34) (cid:1) (cid:27) (cid:34) (cid:6) (cid:34) (cid:4) (cid:2)(cid:34) (cid:8)(cid:24)(cid:26)(cid:31)(cid:13)(cid:26)(cid:15)(cid:34)(cid:28)(cid:17)(cid:24)(cid:14)(cid:19)(cid:34)(cid:10)(cid:16)(cid:30)(cid:16)(cid:26)(cid:28)(cid:16)(cid:34)(cid:28)(cid:17)(cid:24)(cid:14)(cid:19)(cid:34) (cid:3)(cid:2) (cid:1) (cid:8)(cid:4)(cid:2) (cid:3)(cid:3)(cid:3)(cid:3) (cid:1) (cid:3)(cid:3) (cid:1)(cid:2) 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(cid:6)(cid:10)(cid:8)(cid:33)(cid:21)(cid:23)(cid:14)(cid:15)(cid:19)(cid:33)(cid:1)(cid:26)(cid:33)(cid:5)(cid:33)(cid:3)(cid:2)(cid:33)(cid:7)(cid:23)(cid:25)(cid:30)(cid:12)(cid:25)(cid:14)(cid:33)(cid:27)(cid:16)(cid:23)(cid:13)(cid:18)(cid:33)(cid:9)(cid:15)(cid:29)(cid:15)(cid:25)(cid:27)(cid:15)(cid:33)(cid:27)(cid:16)(cid:23)(cid:13)(cid:18)(cid:33) (cid:3)(cid:2) (cid:1) (cid:8)(cid:4)(cid:2) (cid:1)(cid:1) (cid:2)(cid:5) (cid:1)(cid:1)(cid:1) (cid:4) (cid:1)(cid:1)(cid:1) (cid:4) (cid:3)(cid:1) (cid:2)(cid:1) (cid:1) (cid:2)(cid:5) (cid:1) (cid:1)(cid:5)(cid:6) (cid:1)(cid:4) (cid:1) (cid:2)(cid:5) (cid:1)(cid:1) (cid:4)(cid:1)(cid:4) (cid:2)(cid:1) (cid:1)(cid:2) (cid:4)(cid:3) (cid:1)(cid:1) (cid:4) (cid:1)(cid:2) (cid:4)(cid:2) (cid:1)(cid:1) (cid:4) (cid:1)(cid:3)(cid:1)(cid:2) (cid:4)(cid:3) (cid:1)(cid:1) (cid:1) (cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1)(cid:1) (cid:2)(cid:5) (cid:1)(cid:2) (cid:2)(cid:5) 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Figure 10.
Best-fit light curves of SN 2019uo fitted with a CSI model and a combination of Ni and CSI. The forward shocks,reverse shocks, and Ni models are plotted with different lines.
Table 3.
Parameters of the Ni model. The uncertaintiesare 1 σ . M ej M Ni v sc0 κ γ, Ni t (cid:63) expl χ / dof(M (cid:12) ) (M (cid:12) ) (10 cm s − ) (cm g − ) (days)0 . +0 . − . . +0 . − . . +0 . − . . +0 . − . − . +0 . − . . / (cid:63) The value of t expl is with respect to r max . The ejecta can be broadly distinguished into twozones, the inner part ( ρ ej ∝ r − δ ) and the outer part( ρ ej ∝ r − n ). The density profile of the CSM can typ-ically be described as a power law where ρ CSM ∝ r − s ,where s = 0 corresponds to shells of the CSM and s = 2corresponds to winds. Assuming δ = 1 and n = 10, theadopted parameters of the CSM model are the energyof the SN ( E SN ), the mass of the ejecta ( M ej ), the massof the CSM ( M CSM ), the density of the innermost partof the CSM ρ CSM , in , the radius of the innermost partof the ejecta R CSM , in , the efficiency factor which con- verts kinetic energy to radiation ( (cid:15) ), the dimensionless x parameter , and t expl . Two additional parametersare used in the Ni + CSI model, M Ni and κ γ, Ni . Thebest-fit parameters of the model are tabulated in Ta-ble 4 and the best-fit models are displayed in Fig. 10.The corner plots describing covariance of the parame-ters are shown in Fig. 11, Fig. 12, Fig. 13 and Fig. 14respectively. The tabulated values of ejecta masses ofthe four models are reasonable if the progenitor is a WRstar of mass ∼ M (cid:12) and the metallicity is nearly so-lar (Crowther & Smartt 2007). We adopted the Ni,CSI model, and the Ni + CSI models to fit the bolo-metric light curve of SN 2019uo. The Ni model pro-vides a favourable fit to the light curve, but this modelcannot explain the He I emission lines present in thespectrum of SN 2019uo. These lines are likely gener-ated because of the CSI. We therefore invoke CSI as themore favourable model to model light curve. For the x ≡ r ( t ) R ( t ) , where x ≤ x and x ≥ x are inner and outer parts ofthe ejecta. lash ionization signatures in type Ibn supernova SN 2019uo Table 4.
Parameters of the CSI model and the CSI plus Ni model. The uncertainties are 1 σ . s E SN M ej M Ni M CSM ρ CSM , in R CSM , in (cid:15) x κ γ, Ni t (cid:63) expl χ / dof(10 erg) ( M (cid:12) ) ( M (cid:12) ) ( M (cid:12) ) (10 − g cm − ) (10 cm) (cm g − ) (days)CSI 2 0 . +0 . − . . +0 . − . · · · . +0 . − . . +3 . − . . +0 . − . . +0 . − . . +0 . − . · · · − . +0 . − . . / . +0 . − . . +3 . − . · · · . +0 . − . . +0 . − . . +6 . − . . +0 . − . . +0 . − . · · · − . +0 . − . . / Ni 2 1 . +0 . − . . +2 . − . . +0 . − . . +0 . − . . +4 . − . . +1 . − . . +0 . − . . +0 . − . . +10 . − . − . +0 . − . . / Ni 0 1 . +0 . − . . +2 . − . . +0 . − . . +0 . − . . +2 . − . . +1 . − . . +0 . − . . +0 . − . . +10 . − . − . +0 . − . . / (cid:63) The value of t expl is with respect to r max . CSI model, the estimated ejecta masses for s = 0 and s = 2 are 8 . +0 . − . M (cid:12) and 13 . +3 . − . M (cid:12) respectively.This model, however, did not take into account the roleof Ni. Using the combination of both Ni + CSI, theestimated M ej for s = 0 and s = 2 is 15 . +2 . − . M (cid:12) and 16 . +2 . − . M (cid:12) , respectively, which are consistentwith a WR progenitor scenario. The mass-loss rate isgiven by ˙ M = 4 πv w q (where q = ρ CSM , in R , in ). Thevelocity of the wind v w = 100-1000 km s − for WR sys-tems. Considering the wind CSI model (s = 2), we findthat the estimated mass-loss rate lies between 0.195-1.95 M (cid:12) yr − , which is comparable with the values ob-tained for iPTF13z (0.1–2 M (cid:12) yr − ; Nyholm et al. 2017)and PS15dpn (1–10 M (cid:12) yr − ; Wang & Li 2019). Usingthe combination of Ni + CSI model (s = 2), the esti-mated mass loss rate lies between 25.5-255.4 M (cid:12) yr − which is significantly higher than the value obtained foriPTF13z, PS15dpn, and this model can be excluded.Neverthless, the Ni+CSM shell is reasonable.For the CSM shell and the Ni + CSM shell model,the expelled shell mass prior to explosion are ∼ M (cid:12) and 0.73 M (cid:12) , respectively. The radius of the inner shellfor the Ni + CSI model, as seen from Table 4, is 14 × cm and the typical velocity of WR winds is between100 and 1000 km s − (10 − cm s − ); so the time atwhich the shell is expelled prior to explosion is estimatedto be between 1 . × s and 1 . × s, i.e., between163.8 and 1638.8 days. SUMMARYIn this paper, we present the photometric and spec-tral evolution of the type Ibn SN 2019uo. The typicallight curve decay rate of SNe Ibn is ∼ − in allbands which is in agreement with the decline rates of theSNe Ibn discussed by Hosseinzadeh et al. (2017). Thecolor evolution of SN 2019uo is similar to SN 2010al andiPTF14aki which places it between SNe Ib and SNe Ibn. This is in good agreement with the P Cygni spectro-scopic features that transition from narrow to broad,indicating a He-rich circumstellar shell around the pro-genitor star along with optically thick CSM (Hossein-zadeh et al. 2017). The absolute magnitude ( M Vmax = − ± Ni model. However, the Nimodel alone does not take into account the CSM inter-action that is evident from the narrow emission lines inthe spectra of SN 2019uo. Thus, we also fit the lightcurves with a CSI model and a Ni + CSI model. The Ni + CSI wind (s=2) model can be excluded since anunrealistic value of mass loss rate (25.5-255.4 M (cid:12) yr − )is required and the Ni + CSI shell model is reason-able. The combination of Ni + CSI shell well fits ourobserved light curve, with ejecta masses consistent witha WR star. The spectroscopic features of SN 2019uo in-dicate that it is the second SNe Ibn with flash ionizationsignatures. Prominent lines of He II , C III , and N
III aredetected in the spectra, similar to SN 2010al. SN 2019uoshows initial P Cygni He I features that broadens after11 days post-maximum. This can originate from a He-rich shell around progenitor surrounded by dense CSM,or it may be due to viewing angle dependency. Thisis also validated by the equivalent widths of He I fea-tures. Alternatively, P-cygni spectroscopic features usu-ally originate from optical depths ≤
1. As X-rays pene-trate into the P-cygni producing regions absorptions arefilled leading to subsequent emission features. The esti-mated line velocities are lower than the average SN Ibn,but they show a faster evolution compared to the groupof SNe that show prominent emission features from thebeginning. ACKNOWLEDGMENTS2
Gangopadhyay et al.
Figure 11.
The corner plot of the CSI wind model displaying covariance of estimated parameters.
We thank the support of the staff of the Xinglong2.16 m and Lijiang 2.4 m telescope, and observing as-sistants at the 1.04 m ST, 1.30 m DFOT, and 2.00 mHCT for their support during observations. This workwas partially supported by the Open Project Programof the Key Laboratory of Optical Astronomy, NationalAstronomical Observatories, Chinese Academy of Sci-ences. Funding for the Lijiang 2.4 m telescope has been provided by Chinese Academy of Sciences and thePeople’s Government of Yunnan Province. We acknowl-edge Weizmann Interactive Supernova data REPosi-tory (WISeREP; http://wiserep.weizmann.ac.il). DAHacknowledges support from NSF grant AST-1313404.The work of XW is supported supported by the Na-tional Natural Science Foundation of China (NSFCgrants 11325313, 11633002, and 11761141001), and lash ionization signatures in type Ibn supernova SN 2019uo Figure 12.
The corner plot of the CSI shell model displaying covariance of estimated parameters. the National Program on Key Research and Devel-opment Project (grant no. 2016YFA0400803). Thiswork makes use of data obtained with the LCO Net-work. KM and SBP acknowledges BRICS grantDST/IMRCD/BRICS/Pilotcall/ProFCheap/2017(G)for the present work. SBP and KM also acknowledgethe DST/JSPS grant, DST/INT/JSPS/P/281/2018.KM acknowledges the support from Department ofScience and Technology (DST), Govt. of India and Indo-US Science and Technology Forum (IUSSTF)for the WISTEMM fellowship and Dept. of Physics,UC Davis where a part of this work was carriedout. GCA, BK and DKS acknowledge BRICS grantDST/IMRCD/BRICS/PilotCall1/MuMeSTU/2017(G)for the present work. The work of SQW is supported byNational Natural Science Foundation of China (Grants11963001, 11533003, 11603006, 11673006, 11851304,and U1731239).4
Gangopadhyay et al.
Figure 13.
The corner plot of the Ni + CSI wind model displaying covariance of estimated parameters.
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