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SN 2015as and the local standard stars in the field of UGC 5460. R-band, 300 sec image obtained on 2016 February 24 with the 200 cm HCT.
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We present results of the photometric (from 3 to 509 days past explosion) and spectroscopic (up to 230 days past explosion) monitoring campaign of the He-rich Type IIb supernova (SN) 2015as. The {\it (B-V)} colour evolution of SN 2015as closely resemble those of SN 2008ax, suggesting that SN 2015as belongs to the SN IIb subgroup that does not show...
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... Q is the airmass and k λ are the extinction coeffi- cients. To correct for the atmospheric extinction, we used the site extinction values kv = 0.12 ± 0.04, k b = 0.21 ± 0.04, ki = 0.05 ± 0.03 and kr = 0.09 ± 0.04 mag airmass −1 ( Stalin et al. 2008). A root-mean-squared (rms) scatter be- tween transformed and standard magnitude of Landolt stars was found to be between 0.02 -0.03 mag in the BVRI bands. Using these transformation equations, we calibrated 15 non-variable local standards in the SN field. These sec- ondary standards are used to convert the SN instrumental magnitudes into apparent magnitudes in BVRI filters from EKAR/AFOSC, ST/1K x 1K CCD and DFOT/512 x 512 CCD. The local secondary standards in the SN field are marked in Fig 1, and their magnitudes are listed in Table 1. For each night, precise zero points were determined us- ing these secondary standards in order to correct for non- photometric conditions. The errors due to calibration and photometry were added in quadrature to estimate the final error in the SN magnitudes. The final SN magnitudes and their associated errors are listed in Table ...
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... SN magnitudes in the ugriz bands, obtained with the 182cm Ekar Asiago telescope were differentially cali- brated using secondary standards obtained from SDSS DR12 ( Eisenstein et al. 2011). We selected seven out of the 15 stan- dard stars that are marked in Fig 1 with known Sloan ugriz magnitudes. They are listed in Table 2, while the Sloan-band SN magnitudes are listed in Table ...
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... spectra, in the transition from the photo- spheric to the nebular phase, show a weakening of the Ca II H & K feature (see Fig 12). The absorption of Hα decreases in strength with time and, at ∼ 75 days, it marks the meta- morphosis of the SN spectrum from Type II to Type Ib. Lines of He I 4471, 5015, 5876, 6678, 7065 and 7281Å7281Å are now well observed. The presence of Hα and Hβ absorption features until ∼ 75 days indicates that the progenitor was still relatively H-rich at the time of the explosion. A resid- ual Hα emission feature is seen between 75 to 110 days post explosion. The O I line at 7774Å7774Å is detected, in addition to the absorption due to Mg II. The He I 5876Å5876Å wing is becoming prominent and the He I absorption profile shows a notch due to the Na ID feature. This phase is characterised by the appearance of the Ca II NIR emission feature, which gradually becomes more prominent from the 50 to 110 days spectra. The spectra now resemble those of a Type Ib event, except for the residual presence of the Hα feature. The for- bidden line of [O I] 5577Å5577Å , and O I at 7774Å7774Å become more prominent with time. Fig 13 shows that the spectrum of SN 2015as at this phase match well those of SNe 1993J, 2008ax and 2011dh. The Hα profile at this stage is similar to SNe 2003bg and 1996cb, suggesting a similar progenitor configu- ration. The Ca II NIR emission of SN 2015as matches those observed in the spectra of SNe 1993J and 2008ax, although the peaks are more pronounced in SN ...
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... spectra, in the transition from the photo- spheric to the nebular phase, show a weakening of the Ca II H & K feature (see Fig 12). The absorption of Hα decreases in strength with time and, at ∼ 75 days, it marks the meta- morphosis of the SN spectrum from Type II to Type Ib. Lines of He I 4471, 5015, 5876, 6678, 7065 and 7281Å7281Å are now well observed. The presence of Hα and Hβ absorption features until ∼ 75 days indicates that the progenitor was still relatively H-rich at the time of the explosion. A resid- ual Hα emission feature is seen between 75 to 110 days post explosion. The O I line at 7774Å7774Å is detected, in addition to the absorption due to Mg II. The He I 5876Å5876Å wing is becoming prominent and the He I absorption profile shows a notch due to the Na ID feature. This phase is characterised by the appearance of the Ca II NIR emission feature, which gradually becomes more prominent from the 50 to 110 days spectra. The spectra now resemble those of a Type Ib event, except for the residual presence of the Hα feature. The for- bidden line of [O I] 5577Å5577Å , and O I at 7774Å7774Å become more prominent with time. Fig 13 shows that the spectrum of SN 2015as at this phase match well those of SNe 1993J, 2008ax and 2011dh. The Hα profile at this stage is similar to SNe 2003bg and 1996cb, suggesting a similar progenitor configu- ration. The Ca II NIR emission of SN 2015as matches those observed in the spectra of SNe 1993J and 2008ax, although the peaks are more pronounced in SN ...
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... second spectrum was obtained 27 days after the explo- sion, and is shown along with the 29 and 34 day spectra in Fig 11. The most prominent lines, including the He features, are marked in the figure. A common property of the three spectra is the weakening of the blue continuum. We also note a decrease in the strength of Hβ and Hγ while He I features now become prominent. In particular, P Cygni fea- tures due to He I lines 4471, 5015, 5876, 6678, 7065 and 7281Å 7281Å are now strong. The appearance and the strengthening of He I 6678Å6678Å and 7065Å7065Å about a month after the explo- sion, indicates that the progenitor of SN 2015as was partially stripped. The Fe II lines between 4300 and 5000Å5000Å become stronger with time. A double notch in absorption is seen at 5300Å5300Å probably due to Sc II lines. We also see a O I 8448Å 8448Å line near the left edge of Ca II NIR triplet. The O I line later on blends with the increasing Ca II NIR. The Hα absorption line becomes somewhat narrower, and a double trough is seen in 27 and 29 days spectra. The red component is primarily due to the P Cygni profile of Hα line while the origin of the blue component is still not clear. It has been suggested that the blue component is also a part of Hα line, e.g. due to the presence of non-spherical density distribution of H ( Schmidt et al. 1993) or the existence of a second high velocity Hα layer ( Zhang et al. 1995;Branch et al. 2002). The spectral modelling reveals again that a combination of Hα and Si II 6355Å6355Å reproduces the observed features at ...
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... nebular phase spectral evolution from 132 to 230 days after explosion is shown in the Fig 15). Although the He I lines at 6678, 7065Å7065Å are now relatively weaker, the He I 5876Å5876Å is still pronounced in SN 2015as and SN 2011dh, most likely due to an increasing contribution of Na ID. We also estimate the [ has one peak close to the rest velocity while the other one is blueshifted by ∼ -1700 km sec −1 . The possible expla- nations for the observed blueshift are low mode convective instabilities ( Scheck et al. 2004;Kifonidis et al. 2006) or the suppression of the redshifted part of the emission spectrum caused by dust formation as the ejecta cool ( Matheson et al. 2000b;Elmhamdi et al. 2004). A straightforward geometric explanation could be either asymmetric explosions with the emitting oxygen located in a torus or in a disc perpendic- ular to the line of sight ( Mazzali et al. 2001;Maeda et al. 2006b), or a blob of oxygen moving perpendicularly to the line of sight. For different SNe, multiple explanations were proposed for the origin of double peaked profiles. The [O I] line profile, in most cases is the result of clumpy ejecta with a sawtooth profile ( Matheson et al. 2000b), which is supported by the explosion scenario of SN 1987A ( Li et al. 1993) where the O emission originates from clumps of newly synthesised material, while the Ca emission mainly originated from pre- existing, uniformly distributed material. Filippenko & Sargent (1989) suggested that these enhancements come from Rayleigh-Taylor fingers of high speed material or changes in local density contrasts. Asphericity in the explosion would also result in asymmetric peaks ( Milisavljevic et al. 2010). An alternative explanation for the double peaked structure of O lines could be a high-velocity (12,000 km sec Li et al. 1993;Matheson et al. 2000a). Calcium clumps are formed during explosion but they do not contribute significantly to the [Ca II] emis- sion. They intercept γ-ray radiation, however, as the mass fraction is comparatively less, the amount of radioactive lu- minosity and temperature achieved is not sufficient for the ...
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... Fig. 18, we compare the velocity evolution of SN 2015as with those of SNe 1993J, 2003bg, 2008ax, 2010as, 2011dh, 2011ei, 2011fu and 2013df. In the present Type IIb sample, SN 2015as has the lowest observed velocities. At about 10 days after explosion, the Hα velocity for other SNe IIb like 2003bg, 2010as and 2011fu is 17,000 km sec −1 , while SNe 1993J, 2008ax, 2011dh and 2013df is ∼ 13,000 km sec −1 , which are about 17-18 % higher than that of SN 2015as, while SN 2011ei has a similar velocity of 12,500 km sec −1 (see above; and top-left panel of Fig. 18). Even though Hβ and HeI starts with a lower velocity, a similar trend is also observed for the Hβ (top-right in the figure) and the He I lines (bottom-left), while Fe II follows a behaviour similar to other SNe IIb (bottom-right). In SN 2015as, Hα remained visible for a longer time at lower velocities than all the mem- bers of the comparison sample. Iwamoto et al. (1997) showed that the minimum H velocity in SNe IIb mostly depends on the mass of the H envelope retained at the explosion time. leading to lower velocity of the expanding ejecta. Indeed, our observations indicate a peak magnitude towards the faint end of the Type IIb SNe sample distribution, although it is brighter than the faintest SNe 2011ei and ...
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... Fig. 18, we compare the velocity evolution of SN 2015as with those of SNe 1993J, 2003bg, 2008ax, 2010as, 2011dh, 2011ei, 2011fu and 2013df. In the present Type IIb sample, SN 2015as has the lowest observed velocities. At about 10 days after explosion, the Hα velocity for other SNe IIb like 2003bg, 2010as and 2011fu is 17,000 km sec −1 , while SNe 1993J, 2008ax, 2011dh and 2013df is ∼ 13,000 km sec −1 , which are about 17-18 % higher than that of SN 2015as, while SN 2011ei has a similar velocity of 12,500 km sec −1 (see above; and top-left panel of Fig. 18). Even though Hβ and HeI starts with a lower velocity, a similar trend is also observed for the Hβ (top-right in the figure) and the He I lines (bottom-left), while Fe II follows a behaviour similar to other SNe IIb (bottom-right). In SN 2015as, Hα remained visible for a longer time at lower velocities than all the mem- bers of the comparison sample. Iwamoto et al. (1997) showed that the minimum H velocity in SNe IIb mostly depends on the mass of the H envelope retained at the explosion time. leading to lower velocity of the expanding ejecta. Indeed, our observations indicate a peak magnitude towards the faint end of the Type IIb SNe sample distribution, although it is brighter than the faintest SNe 2011ei and ...
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... expansion velocity of the ejecta is usually measured from the absorption minima of the P Cygni profile by fit- ting a gaussian. We choose the relatively isolated lines of Hα, Hβ, He I 5876Å5876Å, Fe II 5169Å5169Å, Ca H & K and the Ca II NIR triplet. We choose 3934Å3934Å as the rest wavelength for Ca II doublet and 8498Å8498Å for the Ca II NIR triplet. The evo- lution of the line velocities is shown in Fig. 17. The Hα line velocity is about 11,300 km sec −1 at 11 days after explosion, and fades to ∼8,000 km sec −1 at 80 days after explosion, be- coming almost constant thereafter. However, the late-time velocity estimates are affected by a large uncertainty, as the blue wing of the absorption profile of Hα is contaminated by the [O I] emission feature. The velocity of Ca II H & K at day 11 is 10,100 km sec −1 , but has a faster decline until ∼30 days, to become almost constant later on. The velocity trend of He I is similar to that of Hα but with comparatively lower velocities. The measured velocity of He I at ∼ 11 days post explosion is 7,000 km sec −1 , drops to 5,500 km sec −1 on day 30 and then remains almost constant. The Ca II NIR feature also has a similar trend starting with a velocity of 8,000 km sec −1 , decreasing fast and becoming constant at 6,500 km sec −1 . Hβ follows a gradual decline with an initial 5000 0 5000 in both photospheric and nebular phases reveal- ing the existence of the stratified layers of weaker Balmer lines at 10,000 km sec −1 followed by He I lines at 7,000 km sec −1 . Similarly, it appears that SN 2015as has a three-layer velocity stratification -an outer layer of H moving with high velocity (∼ 8,000 km sec −1 ); an intermediate Ca-rich layer (7000 km sec −1 ) and a final dense iron core with 4,000 km sec −1 . This scenario can be claimed both from the minima of the absorption profiles that are evolving with time and also from the velocities obtained from the SYN++ modelling. SN 2011dh also showed a similar velocity profile ( Sahu et al. 2013), although the velocities in SN 2011dh were slightly higher than in SN 2015as. It is also important to note that not only stratification but optical depth also plays a vital role in this division of ...
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Optical and near-infrared photometry and optical spectroscopy are reported for SN 2003bg, starting a few days after explosion and extending for a period of more than 300 days. Our early-time spectra reveal the presence of broad, high-velocity Balmer lines. The nebular-phase spectra, on the other hand, show a remarkable resemblance to those of Type...
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... Type IIb SNe are shown by colored hollow rhombus connected by colored solid lines and the other type II SNe are shown by colored solid lines. 2024), SN 2015as ∼ 5.2 × 10 42 erg s −1(Gangopadhyay 2018). And the L peak of SN 2019tua is similar to SN 2020cpg ∼ 6.03 × 10 42 erg s −1(Medler et al. 2021;Teffs et al. 2022). ...
... And the L peak of SN 2019tua is similar to SN 2020cpg ∼ 6.03 × 10 42 erg s −1(Medler et al. 2021;Teffs et al. 2022). We also plotted the LCs from the referenceMedler et al. 2022;Teffs et al. 2022;Gangopadhyay 2018) as shown inFigure 8, with the The bolometric LC of SN 2019tua reproduced by the 56 Ni model. The total luminosity Linp,Ni is represented by the solid blue line. ...
We present photometric and spectroscopic observations and analysis of the type IIb supernova (SN) SN 2019tua, which exhibits multiple bumps in its declining light curves between 40 and 65 days after discovery. SN 2019tua shows a time to peak of about 25 days similar to other type IIb SNe. Our observations indicate a decrease in its brightness of about 1 magnitude in the 60 days after the peak. At about days 50, and 60, its multiband light curves exhibit bumpy behavior. The complex luminosity evolution of SN 2019tua could not be well modeled with a single currently popular energy source model, e.g., radioactive decay of $^{56}$Ni, magnetar, interaction between the ejecta and a circumstellar shell. Even though the magnetar model has a smaller \( \chi^2 / \text{dof} \) value, the complex changes in SN 2019tua's brightness suggest that more than one physical process might be involved. We propose a hybrid CSM interaction plus $^{56}$Ni model to explain the bolometric light curve (LC) of SN 2019tua. The fitting results show that the ejecta mass $M_{\rm ej} \approx 2.4~M_\odot$, the total CSM mass $M_{\rm CSM} \approx 1.0~M_\odot$, and the $^{56}$Ni mass $M_{\rm Ni} \approx 0.4~M_\odot$. The total kinetic energy of the ejecta is $E_k\approx 0.5 \times 10^{51}\rm~erg$. Pre-existing multiple shells suggest that the progenitor of SN 2019tua experienced mass ejections within approximately $\sim6 - 44$ years prior to the explosion.
We present optical, near-infrared, and radio observations of supernova (SN) SN IIb 2022crv. We show that it retained a very thin H envelope and transitioned from an SN IIb to an SN Ib; prominent H α seen in the pre-maximum phase diminishes toward the post-maximum phase, while He i lines show increasing strength. SYNAPPS modeling of the early spectra of SN 2022crv suggests that the absorption feature at 6200 Å is explained by a substantial contribution of H α together with Si ii , as is also supported by the velocity evolution of H α . The light-curve evolution is consistent with the canonical stripped-envelope SN subclass but among the slowest. The light curve lacks the initial cooling phase and shows a bright main peak (peak M V = −17.82 ± 0.17 mag), mostly driven by radioactive decay of ⁵⁶ Ni. The light-curve analysis suggests a thin outer H envelope ( M env ∼ 0.05 M ⊙ ) and a compact progenitor ( R env ∼ 3 R ⊙ ). An interaction-powered synchrotron self-absorption model can reproduce the radio light curves with a mean shock velocity of 0.1 c . The mass-loss rate is estimated to be in the range of (1.9−2.8) × 10 ⁻⁵ M ⊙ yr ⁻¹ for an assumed wind velocity of 1000 km s ⁻¹ , which is on the high end in comparison with other compact SNe IIb/Ib. SN 2022crv fills a previously unoccupied parameter space of a very compact progenitor, representing a beautiful continuity between the compact and extended progenitor scenario of SNe IIb/Ib.
We present a systematic analysis of 191 stripped-envelope supernovae (SE SNe), aimed at computing their ⁵⁶ Ni masses from the luminosity in their radioactive tails ( M Ni tail ) and/or in their maximum light, and the mean ⁵⁶ Ni and iron yields of SE SNe and core-collapse SNe. Our sample consists of SNe IIb, Ib, and Ic from the literature and from the Zwicky Transient Facility Bright Transient Survey. To calculate luminosities from optical photometry, we compute bolometric corrections using 49 SE SNe with optical and near-IR photometry, and develop corrections to account for the unobserved UV and IR flux. We find that the equation of Khatami & Kasen for radioactive ⁵⁶ Ni-powered transients with a single free parameter does not fit the observed peak time–luminosity relation of SE SNe. Instead, we find a correlation between M Ni tail , peak time, peak luminosity, and decline rate, which allows for measuring individual ⁵⁶ Ni masses to a precision of 14%. Applying this method to the whole sample, we find, for SNe IIb, Ib, and Ic, mean ⁵⁶ Ni masses of 0.066 ± 0.006, 0.082 ± 0.009, and 0.132 ± 0.011 M ⊙ , respectively. After accounting for their relative rates, for SE SNe as a whole, we compute mean ⁵⁶ Ni and iron yields of 0.090 ± 0.005 and 0.097 ± 0.007 M ⊙ , respectively. Combining these results with the recent Type II SN mean ⁵⁶ Ni mass derived by Rodríguez et al., core-collapse SNe, as a whole, have mean ⁵⁶ Ni and iron yields of 0.055 ± 0.006 and 0.058 ± 0.007 M ⊙ , respectively. We also find that radioactive ⁵⁶ Ni-powered models typically underestimate the peak luminosity of SE SNe by 60%–70%, suggesting the presence of an additional power source contributing to the luminosity at peak.
This paper studies about the 130-cm Devasthal Fast Optical Telescope (DFOT) at Devasthal, India that has been in operation for more than 10 years and is the main workhorse for the photometric observations for a wide range of scientific programs carried out at ARIES, Nainital. Having a [Formula: see text] pixel imager mounted on the prime focus of the telescope, DFOT provides a field of view of about [Formula: see text] arcmin ² in the sky. Another frame transfer CCD imager of [Formula: see text] pixel size enables monitoring transient sources with millisecond temporal resolution. DFOT is equipped with a filter assembly having eight filters, an auto-guider, an All Sky Camera, and GPS-enabled weather monitoring system to support the observations in the most optimum way. The telescope is capable of producing sub-milimag photometric stability which has allowed us to detect many small-scale photometric variations.
We present a systematic analysis of 191 stripped-envelope supernovae (SE SNe), aimed to compute their $^{56}$Ni masses from the luminosity in their radioactive tails ($M_\mathrm{Ni}^\mathrm{tail}$) and/or in their maximum light, and the mean $^{56}$Ni and iron yields of SE SNe and core-collapse SNe. Our sample consists of SNe of types IIb, Ib, and Ic from the literature and from the Zwicky Transient Facility Bright Transient Survey. We use color curves to infer host galaxy reddenings and the representative $R_V$ value for each SN type. To calculate luminosities from optical photometry, we compute bolometric corrections using 49 SE SNe with optical and near-IR photometry. We find that the equation of Khatami & Kasen relating peak time and luminosity is not a reliable estimator of the $^{56}$Ni masses of SE SNe. Instead, we find a correlation between $M_\mathrm{Ni}^\mathrm{tail}$, peak time, peak luminosity, and decline rate, which allows measuring individual $^{56}$Ni masses to a precision of 14%. Applying this method to the whole sample, we find, for SNe IIb, Ib, and Ic, mean $^{56}$Ni masses of $0.066\pm0.006$, $0.082\pm0.009$, and $0.132\pm0.011$ M$_{\odot}$, respectively. After accounting for their relative rates, for SE SNe as a whole we compute mean $^{56}$Ni and iron yields of $0.090\pm0.005$ and $0.097\pm0.006$ M$_{\odot}$, respectively. Combining these results with the recent Type II SN mean $^{56}$Ni mass derived by Rodr\'iguez et al., core-collapse SNe, as a whole, have mean $^{56}$Ni and iron yields of $0.055\pm0.006$ and $0.058\pm0.007$ M$_{\odot}$, respectively. We highlight that the Arnett model, Arnett's rule, and hydrodynamical models typically overestimate the $^{56}$Ni masses of SE SNe by 75, 90 and 65%, respectively.
We report results of optical imaging and low-resolution spectroscopic monitoring of supernova (SN) 2017iro that occurred in the nearby (∼31 Mpc) galaxy NGC 5480. The He i λ 5876 feature present in the earliest spectrum (−7 days) classified it as a Type Ib SN. The follow-up observations span from −7 to +266 days with respect to the B -band maximum. With a peak absolute magnitude in V band M V = −17.76 ± 0.15 mag and bolometric luminosity log 10 L = 42.39 ± 0.09 erg s ⁻¹ , SN 2017iro is a moderately luminous Type Ib SN. The overall light-curve evolution of SN 2017iro is similar to that of SN 2012au and SN 2009jf during the early (up to ∼100 days) and late phases (>150 days), respectively. The line velocities of both Fe ii λ 5169 and He i λ 5876 are ∼9000 km s ⁻¹ near the peak. The analysis of the nebular phase spectrum (∼+209 days) indicates an oxygen mass of ∼0.35 M ⊙ . The smaller [O i ]/[Ca ii ] flux ratio of ∼1 favors a progenitor with a zero-age main-sequence mass in the range ∼13–15 M ⊙ , most likely in a binary system, similar to the case of iPTF13bvn. The explosion parameters are estimated by applying different analytical models to the quasi-bolometric light curve of SN 2017iro. ⁵⁶ Ni mass synthesized in the explosion has a range of ∼0.05–0.10 M ⊙ , ejecta mass ∼1.4–4.3 M ⊙ , and kinetic energy ∼(0.8–1.9) × 10 ⁵¹ erg.
We report results of optical imaging and low-resolution spectroscopic monitoring of supernova (SN) 2017iro that occurred in the nearby ($\sim$\,31 Mpc) galaxy NGC 5480. The \ion{He}{1} 5876 \AA\, feature present in the earliest spectrum (--\,7 d) classified it as a Type Ib SN. The follow-up observations span from --\,7 to +\,266 d with respect to the $B$-band maximum. With a peak absolute magnitude in $V$-band, ($M_{V}$)\,=\,$-17.76\pm0.15$ mag and bolometric luminosity (log$_{10}$\,L)\,=\,42.39\,$\pm$\, 0.09 erg s$^{-1}$, SN 2017iro is a moderately luminous Type Ib SN. The overall light curve evolution of SN 2017iro is similar to SN 2012au and SN 2009jf during the early (up to $\sim$100 d) and late phases ($>$150 d), respectively. The line velocities of both \ion{Fe}{2} 5169 \AA\, and \ion{He}{1} 5876 \AA\, are $\sim$\,9000 km s$^{-1}$ near the peak. The analysis of the nebular phase spectrum ($\sim$\,+209 d) indicates an oxygen mass of $\sim$\,0.35 M$_{\odot}$. The smaller [\ion{O}{1}]/[\ion{Ca}{2}] flux ratio of $\sim$\,1 favours a progenitor with a zero-age main-sequence mass in the range $\sim$\,13--15 M$_{\odot}$, most likely in a binary system, similar to the case of iPTF13bvn. The explosion parameters are estimated by applying different analytical models to the quasi-bolometric light curve of SN 2017iro. $^{56}$Ni mass synthesized in the explosion has a range of $\sim$\,0.05\,--\,0.10 M$_{\odot}$, the ejecta mass $\sim$1.4\,--\,4.3 M$_{\odot}$ and the kinetic energy $\sim$\,0.8\,--\,1.9\,$\times$ 10$^{51}$ erg.
We present the photometric and spectroscopic evolution of the Type II supernova (SN II) SN 2017ivv (also known as ASASSN-17qp). Located in an extremely faint galaxy (Mr = −10.3 mag), SN 2017ivv shows an unprecedented evolution during the 2 yr of observations. At early times, the light curve shows a fast rise (∼6−8 d) to a peak of ${\it M}^{\rm max}_{g}= -17.84$ mag, followed by a very rapid decline of 7.94 ± 0.48 mag per 100 d in the V band. The extensive photometric coverage at late phases shows that the radioactive tail has two slopes, one steeper than that expected from the decay of 56Co (between 100 and 350 d), and another slower (after 450 d), probably produced by an additional energy source. From the bolometric light curve, we estimated that the amount of ejected 56Ni is ∼0.059 ± 0.003 M⊙. The nebular spectra of SN 2017ivv show a remarkable transformation that allows the evolution to be split into three phases: (1) Hα strong phase (<200 d); (2) Hα weak phase (between 200 and 350 d); and (3) Hα broad phase (>500 d). We find that the nebular analysis favours a binary progenitor and an asymmetric explosion. Finally, comparing the nebular spectra of SN 2017ivv to models suggests a progenitor with a zero-age main-sequence mass of 15–17 M⊙.
We present an extensive (∼1200 d) photometric and spectroscopic monitoring of the Type IIn supernova (SN) 2012ab. After a rapid initial rise leading to a bright maximum (MR = −19.39 mag), the light curves show a plateau lasting about 2 months followed by a steep decline up to about 100 d. Only in the U band, the decline is constant in the same interval. At later phases, the light curves remain flatter than the 56Co decline, suggesting the increasing contribution of the interaction between SN ejecta with circumstellar material (CSM). Although heavily contaminated by emission lines of the host galaxy, the early spectral sequence (until 32 d) shows persistent narrow emissions, indicative of slow unshocked CSM, and the emergence of broad Balmer lines of hydrogen with P-Cygni profiles over a blue continuum, arising from a fast expanding SN ejecta. From about 2 months to ∼1200 d, the P-Cygni profiles are overcome by intermediate width emissions [full width at half-maximum (FWHM) ∼6000 km s−1], produced in the shocked region due to interaction. On the red wing, a red bump appears after 76 d, likely a signature of the onset of interaction of the receding ejecta with the CSM. The presence of fast material both approaching and then receding is suggestive that we are observing the SN along the axis of a jet-like ejection in a cavity devoid of or uninterrupted by CSM in the innermost regions.
Type IIb supernovae (SNe IIb) present a unique opportunity for investigating the evolutionary channels and mechanisms governing the evolution of stripped-envelope SN progenitors due to a variety of observational constraints. Comparison of these constraints with the full distribution of theoretical properties not only helps determine the prevalence of observed properties in nature, but can also reveal currently unobserved populations. In this follow-up paper, we use the large grid of models presented in Sravan et al. to derive distributions of single and binary SNe IIb progenitor properties and compare them to constraints from three independent observational probes: multiband SN light curves, direct progenitor detections, and X-ray/radio observations. Consistent with previous work, we find that while current observations exclude single stars as SN IIb progenitors, SN IIb progenitors in binaries can account for them. We also find that the distributions indicate the existence of an unobserved dominant population of binary SNe IIb at low metallicity that arise due to mass transfer initiated on the Hertzsprung Gap. In particular, our models indicate the existence of a group of highly stripped (envelope mass ∼0.1–0.2 M ☉ ) progenitors that are compact (<50 R ☉ ) and blue ( T eff ≲ 10 ⁵ K) with ∼10 4.5 –10 5.5 L ☉ and low-density circumstellar mediums. As discussed in Sravan et al., this group is necessary to account for SN IIb fractions and likely exist regardless of metallicity. The detection of the unobserved populations indicated by our models would support weak stellar winds and inefficient mass transfer in SN IIb progenitors.
