# Investigating the observational properties of Type Ib supernova SN 2017iro

Brajesh Kumar Aryabhatta Research Institute of Observational Sciences, Manora
Peak, Nainital - 263001, India  Indian Institute of Astrophysics, II Block,
Koramangala, Bengaluru 560034, India  Avinash Singh Hiroshima Astrophysical
Science Center, Hiroshima University, Higashi-Hiroshima, Hiroshima 739-8526,
Japan  Indian Institute of Astrophysics, II Block, Koramangala, Bengaluru
560034, India  D.K. Sahu Indian Institute of Astrophysics, II Block,
Koramangala, Bengaluru 560034, India  G.C. Anupama Indian Institute of
Astrophysics, II Block, Koramangala, Bengaluru 560034, India

(Received –; Revised –; Accepted –)

###### Abstract

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 He I 5876 Å 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, ($M_{V}$) = $-17.76\pm 0.15$ mag and bolometric
luminosity (log10 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 Fe II 5169 Å
and He I 5876 Å 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⊙. The smaller [O I]/[Ca II] flux ratio of $\sim$ 1 favours a progenitor with
a zero-age main-sequence mass in the range $\sim$ 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. 56Ni mass synthesized in the explosion has a range
of $\sim$ 0.05 – 0.10 M⊙, the ejecta mass $\sim$1.4 – 4.3 M⊙ and the kinetic
energy $\sim$ 0.8 – 1.9 $\times$ 1051 erg.

supernovae: general – supernovae: individual: SN 2017iro, galaxies:
individual: NGC 5480

††journal: ApJ

##  1 Introduction

Type Ib, Ic, Ic-broad lined (Ic-BL), and IIb supernovae (SNe) belong to the
class of stripped-envelope supernovae (SE-SNe, Minkowski, 1941; Clocchiatti &
Wheeler, 1997; Filippenko, 1997; Gal-Yam, 2017; Modjaz et al., 2019). SE-SNe
are generally characterized by bell-shaped light curves (LCs) which are mainly
powered by the thermal energy released as a consequence of radioactive decay
of the iron group elements (56Ni $\rightarrow$ 56Co $\rightarrow$ 56Fe)
synthesized in the explosion. SE-SNe show no hydrogen (Ib) or neither hydrogen
nor helium (Ic) in their spectra except for Type IIb events, which exhibit
hydrogen in the early phase spectra (Filippenko, 1988; Filippenko, Matheson, &
Ho, 1993). This indicates that they are the results of explosion of stars that
are devoid of the outer layer of hydrogen (Ib) or both hydrogen and helium
(Ic) at the time of explosion. The spectroscopic behaviour of SE-SNe separates
them from Type II SNe, where hydrogen is found to be dominant in the spectra.
The SE-SNe and Type II events are collectively grouped as core-collapse
supernovae (CC-SNe) as they result from the gravitational collapse of the iron
core of massive stars ($M\geq$ 8 M☉) during the end stages of their lives
(Heger et al., 2003; Smartt, 2009).

In SE-SNe, the mechanisms responsible for removing the outer envelopes
(hydrogen or helium) of the progenitors can be either radiation-driven stellar
winds (Conti, 1975; Puls, Vink, & Najarro, 2008; Pauldrach, Vanbeveren, &
Hoffmann, 2012), eruptive mass loss (Smith & Owocki, 2006; Smith, 2006) and/or
mass transfer to a companion star (Nomoto, 1984, 1987; Wheeler & Levreault,
1985; Podsiadlowski, Joss, & Hsu, 1992; Woosley & Weaver, 1995; Nomoto,
Iwamoto, & Suzuki, 1995; Wellstein & Langer, 1999; Wellstein, Langer, & Braun,
2001; Pols & Dewi, 2002; Podsiadlowski, Langer, Poelarends, Rappaport, Heger,
& Pfahl, 2004; Yoon, Woosley, & Langer, 2010). Depending on the main-sequence
mass of the progenitor star, above mechanisms or a combination of them
operates. In the case of a single massive ($>$ 20 M☉) Wolf-Rayet star, the
radiation-driven stellar wind is the dominant process (see Woosley et al.,
1995; Massey, 2003; Crowther, 2007; Smith, 2014a; Yoon, 2015a, and references
therein). In case of moderately massive (8 – 25 M☉) stars, mass transfer to
the binary companion effectively removes the outer envelope (Groh, Georgy, &
Ekström, 2013a, b; Smith, 2014b; Yoon, 2015b, and references therein).

Dedicated and large area transient survey programmes have made it possible to
discover a large number of SNe of different types. The SE-SNe constitute a
reasonable fraction ($\sim$25 – 30%) of the overall SN rate in the local
Universe (Li et al., 2011a; Eldridge et al., 2013; Graur et al., 2017a, b).
The availability of ample data has led to a qualitative analysis of a large
sample of SE-SNe in the recent years (see, e.g. Drout et al., 2011; Bianco et
al., 2014; Modjaz et al., 2014a; Taddia et al., 2015; Liu et al., 2016; Lyman
et al., 2016; Prentice et al., 2016; Taddia et al., 2018; Prentice et al.,
2019; Shivvers et al., 2019; Williamson et al., 2019; Shahbandeh et al.,
2021). These studies reveal that heterogeneity exists in the observational
properties of SE-SNe in terms of synthesized 56Ni masses ($\sim$0.1 – 0.4 M☉),
ejecta masses ($\sim$1 – 6 M☉), explosion energies (few 1051 erg), and
absolute magnitude ($\sim$–17 to –18 mag). Notably, the sample selection
criteria and the methods applied may induce inconsistency in various estimates
(see, Taddia et al., 2018; Prentice et al., 2019).

Statistical investigations of a large sample of photometric data are useful in
exploring the bulk properties of various types of events. Nevertheless, such
studies are generally limited to LCs in selected pass-bands, and also, the
follow-up covers a short duration of the SN evolution. Further, the
spectroscopic follow-up of most objects is restricted to early phases. Non-
uniformity in the data sample may also be present as these are collected at
different observing facilities and detectors. Studying individual events with
proper monitoring at different phases (both photospheric and nebular) is hence
extremely important. The very early phase observations (hours to days after
explosion) of these events are useful to constrain the progenitor radius at
its end stage. This needs a very early detection and quick follow up which is
not always possible considering their random occurrence in the sky. The large
area surveys with high cadence have contributed significantly in this regard.
SN 1993J (Type IIb) is the first event among SE-SNe which shows evidence of a
prominent cooling tail just after explosion (Richmond et al., 1994; Barbon et
al., 1995). During last two decades several Type IIb SNe with such interesting
features have been monitored and studied well e.g. SN 2008ax (Pastorello et
al., 2008; Roming et al., 2009), SN 2011dh (Arcavi et al., 2011), SN 2011fu
(Kumar et al., 2013; Morales-Garoffolo et al., 2015), SN 2013df (Fremling et
al., 2014) and SN 2016gkg (Arcavi et al., 2017; Tartaglia et al., 2017a). A
handful of Type Ib events have also been discovered at very early phases and
their observational properties are studied in detail such as SN 1999ex (Hamuy
et al., 2002; Stritzinger et al., 2002), SN 2008D (Soderberg et al., 2008;
Mazzali et al., 2008; Malesani et al., 2009; Modjaz et al., 2009a; Bersten et
al., 2013), iPTF13bvn (Fremling et al., 2016; Bersten et al., 2014a) and
LSQ13abf (Stritzinger et al., 2020). Recently, Prentice et al. (2019) analyzed
the properties of 18 SE-SNe and discussed the implications for their
progenitors. Direct detection of SE-SNe progenitor candidates has been
possible only for a few objects, e.g. SN 1993J (Maund et al., 2004), SN 2011dh
(Maund et al., 2011; Van Dyk et al., 2013), SN 2001ig (Ryder et al., 2018),
iPTF13bvn (Cao et al., 2013; Eldridge et al., 2015a; Kuncarayakti et al.,
2015a; Folatelli et al., 2016), SN 2016gkg (Kilpatrick et al., 2017; Tartaglia
et al., 2017a), SN 2017ein (Kilpatrick et al., 2018; Van Dyk et al., 2018;
Xiang et al., 2019) and SN 2019yvr (Kilpatrick et al., 2021). Along with very
early phases, the temporal observations during maximum to nebular phases are
equally important to estimate various explosion parameters and progenitor
properties. Detail investigation of more events can provide an alternative way
to understand various progenitor channels.

SN 2017iro was discovered by Patrick Wiggins on 2017 November 30.5 UT
(Wiggins, 2017) at about 15 arcsec South-East of the centre of the galaxy NGC
5480 which also hosted another Type Ib event SN 1988L (Perlmutter et al.,
1988). Pre-explosion images of host galaxy, obtained by Hubble Space Telescope
(HST) in F606W-band using Wide Field Camera (WFC) of Advanced Camera for
Surveys (ACS), exist in the HST archive. SN 2017iro was discovered at a
magnitude of $\sim$15.5 mag (clear filter) and classified as a Type Ib/c SN
with the help of the spectrum taken on 2017 December 02 (Bertrand, 2017). To
be noted that SN 2017iro discovery was slightly late ($\sim$ 8 d after
explosion) however, we were able to cover its peak as well as the late phase
observations with a good sampling of photometric and spectroscopic data
points. Table 1 lists details of SN 2017iro and its host galaxy.

Results based on well-sampled optical photometric and low-resolution
spectroscopic observations of SN 2017iro are presented in this paper. The
observations and data reduction are described in Section 2. Light curve and
spectral properties are presented in Section 3 and 4, respectively. In Section
5, the results are discussed, and the summary is provided in Section 6.

Figure 1: Identification chart of SN 2017iro obtained in $V$-band (SN is
marked with blue crosshairs). The field of view is roughly 9 arcmin $\times$ 9
arcmin. The IDs 1–14 indicate secondary standard stars in the SN field (their
magnitudes are listed in Table 2). The location of SN 1988L is also marked
with a bold circle (yellow). Table 1: Details of SN 2017iro and its host
galaxy.

Parameters | Value | Reference
---|---|---
SN 2017iro: |  |
RA (J2000) |  14h 06m 23$\fs$11 | 1
DEC (J2000) |  +50° 43′ 20$\farcs$20 | 1
Explosion epoch (UT) | 2017 November 23 |  Section 3.1
| (JD 2458080.7) |
Discovery date (UT) | 2017 November 30.5 | 1
| (JD 2458088.0) |
$E(B-V)_{total}$ |  0.28 $\pm$ 0.04 mag |  Section 3.3
NGC 5480: |  |
Type | SA(s)c | 2
Redshift |  0.006191 $\pm$ 0.000017 | 2
Distance |  30.8 $\pm$ 2.2 Mpc | 3
Distance modulus |  32.44 $\pm$ 0.15 mag | 3

1 https://wis-tns.weizmann.ac.il/object/2017iro

2 NASA/IPAC Extragalactic Database (NED)

3 Mould et al. (2000)

Table 2: Calibrated magnitudes of secondary standards in the field of SN 2017iro. Star IDs are indicated in Fig. 1. Star | $U$ | $B$ | $V$ | $R$ | $I$
---|---|---|---|---|---
ID | (mag) | (mag) | (mag) | (mag) | (mag)
1 |  15.86 $\pm$ 0.07 |  15.93 $\pm$ 0.02 |  15.34 $\pm$ 0.01 |  14.95 $\pm$ 0.02 |  14.56 $\pm$ 0.01
2 |  16.16 $\pm$ 0.06 |  16.39 $\pm$ 0.01 |  15.87 $\pm$ 0.01 |  15.53 $\pm$ 0.02 |  15.20 $\pm$ 0.01
3 |  18.02 $\pm$ 0.02 |  17.21 $\pm$ 0.02 |  16.21 $\pm$ 0.01 |  15.59 $\pm$ 0.01 |  15.07 $\pm$ 0.01
4 |  18.09 $\pm$ 0.09 |  17.26 $\pm$ 0.01 |  16.22 $\pm$ 0.01 |  15.57 $\pm$ 0.02 |  15.00 $\pm$ 0.01
5 |  17.21 $\pm$ 0.06 |  17.10 $\pm$ 0.01 |  16.41 $\pm$ 0.01 |  15.97 $\pm$ 0.02 |  15.54 $\pm$ 0.01
6 |  19.83 $\pm$ 0.05 |  19.33 $\pm$ 0.05 |  17.85 $\pm$ 0.02 |  16.50 $\pm$ 0.02 |  14.91 $\pm$ 0.02
7 |  17.62 $\pm$ 0.01 |  17.74 $\pm$ 0.01 |  17.14 $\pm$ 0.01 |  16.76 $\pm$ 0.02 |  16.41 $\pm$ 0.02
8 |  17.33 $\pm$ 0.04 |  17.51 $\pm$ 0.02 |  16.95 $\pm$ 0.02 |  16.59 $\pm$ 0.03 |  16.24 $\pm$ 0.03
9 |  18.14 $\pm$ 0.02 |  17.96 $\pm$ 0.01 |  17.23 $\pm$ 0.01 |  16.78 $\pm$ 0.02 |  16.38 $\pm$ 0.01
10 |  19.92 $\pm$ 0.69 |  18.90 $\pm$ 0.04 |  17.48 $\pm$ 0.02 |  16.30 $\pm$ 0.02 |  14.92 $\pm$ 0.01
11 |  18.27 $\pm$ 0.01 |  18.33 $\pm$ 0.02 |  17.68 $\pm$ 0.02 |  17.28 $\pm$ 0.02 |  16.84 $\pm$ 0.06
12 |  18.17 $\pm$ 0.05 |  17.36 $\pm$ 0.01 |  16.40 $\pm$ 0.01 |  15.80 $\pm$ 0.02 |  15.29 $\pm$ 0.01
13 |  18.04 $\pm$ 0.01 |  18.11 $\pm$ 0.01 |  17.51 $\pm$ 0.01 |  17.12 $\pm$ 0.02 |  16.74 $\pm$ 0.01
14 |  18.66 $\pm$ 0.02 |  18.52 $\pm$ 0.02 |  17.80 $\pm$ 0.02 |  17.36 $\pm$ 0.03 |  16.95 $\pm$ 0.03

##  2 Observations and data reduction

###  2.1 Photometric observation

The photometric follow-up of SN 2017iro started just after the announcement of
its discovery, i.e., 2017 December 01 (JD 2458089.5), and continued up to 2018
August 31 (JD 2458362.1). These observations were performed in Bessell $UBVRI$
bands with the Himalayan Faint Object Spectrograph Camera (HFOSC), mounted on
the $f/9$ Cassegrain focus of the 2 m Himalayan Chandra Telescope (HCT) of
Indian Astronomical Observatory (IAO), Hanle, India (Prabhu, 2014). HFOSC is
equipped with a liquid nitrogen-cooled 2k $\times$ 4k pixels SITe CCD chip
(pixel size 15 $\times$ 15 $\mu$m). The gain and readout noise of the detector
are 1.22 e-/ADU and 4.87 e-, respectively. With a plate scale of 0.296 arcsec
per pixel, the central 2k $\times$ 2k region covers a field of 10$\arcmin$
$\times$ 10$\arcmin$ on the sky and is used for imaging. Further description
on the HCT and the HFOSC instrument can be obtained from
https://www.iiap.res.in/?q=iao_about.

During the observations, several bias and twilight flat frames were obtained
along with science frames. The usual pre-processing steps, such as bias-
subtraction, flat-fielding correction, and cosmic ray removal, were applied to
raw images of the SN. For this purpose, standard tasks available in the data
reduction software IRAF111IRAF stands for Image Reduction and Analysis
Facility. It is distributed by the National Optical Astronomy Observatory,
which is operated by the Association of Universities for Research in Astronomy
(AURA) under a cooperative agreement with the National Science Foundation.
were used. In order to achieve a better signal-to-noise ratio, multiple frames
were taken on some nights and co-added in respective bands after the alignment
of the images.

Table 3: Log of photometric observations and estimated $UBVRI$ magnitudes of
SN 2017iro.

Date | JD |  Phase∗ | $U$ | $B$ | $V$ | $R$ | $I$
---|---|---|---|---|---|---|---
(yyyy-mm-dd) | 2450000+ | (d) | (mag) | (mag) | (mag) | (mag) | (mag)
2017-12-01 | 8089.5 | – 6.7 |  15.85 $\pm$ 0.11 |  16.29 $\pm$ 0.03 |  16.01 $\pm$ 0.11 |  15.78 $\pm$ 0.04 |  15.67 $\pm$ 0.07
2017-12-02 | 8090.5 | – 5.7 |  15.71 $\pm$ 0.11 |  16.20 $\pm$ 0.13 |  15.93 $\pm$ 0.11 |  15.71 $\pm$ 0.03 |  15.56 $\pm$ 0.08
2017-12-05 | 8093.4 | – 2.8 |  15.64 $\pm$ 0.11 |  15.96 $\pm$ 0.13 |  15.68 $\pm$ 0.05 |  15.47 $\pm$ 0.09 |  15.33 $\pm$ 0.11
2017-12-06 | 8094.4 | – 1.8 |  15.64 $\pm$ 0.11 |  16.01 $\pm$ 0.09 |  15.63 $\pm$ 0.11 |  15.38 $\pm$ 0.05 |  15.21 $\pm$ 0.07
2017-12-07 | 8095.5 | – 0.7 |  15.63 $\pm$ 0.12 |  15.95 $\pm$ 0.05 |  15.56 $\pm$ 0.03 |  15.25 $\pm$ 0.02 |  15.118 $\pm$ 0.03
2017-12-08 | 8096.5 | \+ 0.3 |  15.68 $\pm$ 0.12 |  15.92 $\pm$ 0.06 |  15.56 $\pm$ 0.10 |  15.22 $\pm$ 0.10 |  15.10 $\pm$ 0.13
2017-12-09 | 8097.5 | \+ 1.3 | – | – |  15.52 $\pm$ 0.02 |  15.20 $\pm$ 0.03 |  15.06 $\pm$ 0.03
2017-12-12 | 8100.5 | \+ 4.3 |  15.84 $\pm$ 0.12 |  16.12 $\pm$ 0.11 |  15.58 $\pm$ 0.01 |  15.19 $\pm$ 0.04 |  14.99 $\pm$ 0.06
2017-12-14 | 8102.5 | \+ 6.3 | – |  16.33 $\pm$ 0.05 |  15.69 $\pm$ 0.04 |  15.19 $\pm$ 0.02 |  14.96 $\pm$ 0.08
2017-12-17 | 8105.4 | \+ 9.2 |  17.00 $\pm$ 0.13 |  16.74 $\pm$ 0.03 |  15.84 $\pm$ 0.05 |  15.32 $\pm$ 0.02 |  15.06 $\pm$ 0.04
2017-12-26 | 8114.4 | \+ 18.2 |  18.50 $\pm$ 0.14 |  17.76 $\pm$ 0.04 |  16.43 $\pm$ 0.02 |  15.81 $\pm$ 0.04 |  15.38 $\pm$ 0.04
2017-12-29 | 8117.4 | \+ 21.2 | – | – |  16.61 $\pm$ 0.02 |  15.98 $\pm$ 0.02 |  15.59 $\pm$ 0.02
2018-01-05 | 8124.5 | \+ 28.3 | – |  18.28 $\pm$ 0.04 |  16.90 $\pm$ 0.06 |  16.27 $\pm$ 0.02 |  15.77 $\pm$ 0.07
2018-01-07 | 8126.4 | \+ 30.2 |  19.07 $\pm$ 0.14 | – |  16.96 $\pm$ 0.10 |  16.32 $\pm$ 0.08 |  15.84 $\pm$ 0.07
2018-01-08 | 8127.4 | \+ 31.2 | – | – |  16.96 $\pm$ 0.06 |  16.25 $\pm$ 0.11 |  15.79 $\pm$ 0.09
2018-01-14 | 8133.4 | \+ 37.2 | – |  18.40 $\pm$ 0.03 |  17.14 $\pm$ 0.01 |  16.51 $\pm$ 0.02 |  15.99 $\pm$ 0.03
2018-01-16 | 8135.4 | \+ 39.2 |  19.10 $\pm$ 0.20 |  18.50 $\pm$ 0.04 |  17.21 $\pm$ 0.02 |  16.57 $\pm$ 0.04 |  16.03 $\pm$ 0.06
2018-01-17 | 8136.5 | \+ 40.3 | – |  18.52 $\pm$ 0.06 |  17.24 $\pm$ 0.02 |  16.62 $\pm$ 0.02 |  16.05 $\pm$ 0.05
2018-01-21 | 8140.4 | \+ 44.2 | – |  18.49 $\pm$ 0.03 |  17.31 $\pm$ 0.02 |  16.72 $\pm$ 0.01 |  16.12 $\pm$ 0.05
2018-01-30 | 8149.3 | \+ 53.1 | – |  18.66 $\pm$ 0.02 |  17.40 $\pm$ 0.10 |  16.86 $\pm$ 0.10 |  16.19 $\pm$ 0.04
2018-02-02 | 8152.4 | \+ 56.2 | – |  18.57 $\pm$ 0.05 |  17.50 $\pm$ 0.05 |  16.94 $\pm$ 0.03 |  16.34 $\pm$ 0.04
2018-02-03 | 8153.3 | \+ 57.1 |  19.22 $\pm$ 0.21 |  18.61 $\pm$ 0.12 |  17.51 $\pm$ 0.07 |  16.94 $\pm$ 0.06 |  16.32 $\pm$ 0.07
2018-02-09 | 8159.4 | \+ 63.2 | – |  18.65 $\pm$ 0.02 |  17.63 $\pm$ 0.02 |  16.98 $\pm$ 0.02 |  16.41 $\pm$ 0.04
2018-02-13 | 8163.3 | \+ 67.1 |  19.08 $\pm$ 0.15 |  18.73 $\pm$ 0.03 |  17.69 $\pm$ 0.02 |  17.12 $\pm$ 0.03 |  16.50 $\pm$ 0.04
2018-02-16 | 8166.3 | \+ 70.1 |  19.14 $\pm$ 0.17 |  18.70 $\pm$ 0.07 |  17.76 $\pm$ 0.02 |  17.20 $\pm$ 0.03 |  16.57 $\pm$ 0.05
2018-02-20 | 8170.3 | \+ 74.1 | – |  18.77 $\pm$ 0.02 |  17.83 $\pm$ 0.01 |  17.28 $\pm$ 0.02 |  16.64 $\pm$ 0.04
2018-02-25 | 8175.2 | \+ 79.0 | – |  18.79 $\pm$ 0.11 |  17.90 $\pm$ 0.03 |  17.36 $\pm$ 0.03 |  16.72 $\pm$ 0.03
2018-03-07 | 8185.2 | \+ 89.0 | – |  18.96 $\pm$ 0.07 |  18.09 $\pm$ 0.02 |  17.53 $\pm$ 0.02 |  16.87 $\pm$ 0.05
2018-03-15 | 8193.3 | \+ 97.1 | – | – |  18.22 $\pm$ 0.01 |  17.69 $\pm$ 0.01 |  17.05 $\pm$ 0.02
2018-03-23 | 8201.4 | \+ 105.2 | – |  19.18 $\pm$ 0.02 |  18.35 $\pm$ 0.02 |  17.86 $\pm$ 0.02 |  17.19 $\pm$ 0.03
2018-03-30 | 8208.3 | \+ 112.1 | – | – |  18.47 $\pm$ 0.05 |  17.92 $\pm$ 0.04 |  17.30 $\pm$ 0.03
2018-04-02 | 8211.4 | \+ 115.2 | – |  19.28 $\pm$ 0.09 |  18.50 $\pm$ 0.05 |  18.02 $\pm$ 0.03 |  17.38 $\pm$ 0.05
2018-04-07 | 8216.3 | \+ 120.1 | – | – |  18.56 $\pm$ 0.02 |  17.99 $\pm$ 0.02 |  17.38 $\pm$ 0.02
2018-04-15 | 8224.3 | \+ 128.1 | – |  19.50 $\pm$ 0.02 |  18.71 $\pm$ 0.01 |  18.20 $\pm$ 0.01 |  17.60 $\pm$ 0.03
2018-04-21 | 8230.2 | \+ 134.0 | – |  19.64 $\pm$ 0.02 |  18.85 $\pm$ 0.03 |  18.30 $\pm$ 0.01 |  17.76 $\pm$ 0.02
2018-04-28 | 8237.4 | \+ 141.2 | – |  19.79 $\pm$ 0.17 |  18.98 $\pm$ 0.06 |  18.34 $\pm$ 0.04 |  17.86 $\pm$ 0.04
2018-05-01 | 8240.3 | \+ 144.1 | – | – |  19.03 $\pm$ 0.04 |  18.47 $\pm$ 0.03 |  17.88 $\pm$ 0.04
2018-05-30 | 8269.2 | \+ 173.0 | – | – | – |  18.83 $\pm$ 0.02 |  18.38 $\pm$ 0.04
2018-07-05 | 8305.2 | \+ 209.0 | – | – |  20.20 $\pm$ 0.03 |  19.41 $\pm$ 0.02 |  18.96 $\pm$ 0.03
2018-08-31 | 8362.1 | \+ 265.9 | – | – |  21.22 $\pm$ 0.14 |  20.28 $\pm$ 0.12 |  20.16 $\pm$ 0.15

∗ With reference to the $B$-band maximum (JD 2458096.2).

To calibrate a sequence of secondary standards in the SN field, we observed
Landolt photometric standard fields (Landolt, 1992) PG0231+051, PG0918+029,
PG0942-029, and PG1047+003 on three nights, 2018, January 21, February 03 and
25, under good photometric conditions along with the SN field. The observed
Landolt field stars cover a brightness range of 12.27 $\leq V\leq$ 16.11 mag
and colour range of $-0.33\leq B-V\leq 1.45$ mag. The average atmospheric
extinction values in $U$, $B$, $V$, $R$ and $I$ bands for the site were
adopted from Stalin et al. (2008). The average colour terms for the telescope
detector system were used to determine the photometric zero-points by applying
a linear relationship between the observed and standard colours. A set of 14
stars in the SN field (marked in Fig. 1), were calibrated using the estimated
zero-points and average colour terms. The calibrated $UBVRI$ magnitudes of the
secondary standards averaged over three nights are listed in Table 2.
Considering the proximity of SN to an H II region, template subtraction was
performed using the template images of the host galaxy, observed on 2019
August 21 with the HCT. The instrumental magnitudes of the supernova at
different epochs were extracted from the template subtracted images. The SN
magnitude is calibrated differentially with respect to the secondary standards
in the SN field by applying the nightly zero points (see Kumar et al., 2018).
The photometric magnitudes of SN along with the errors are listed in Table 3.
Here, photometric errors include the errors in aperture photometry as
estimated in IRAF and uncertainties in the nightly zero points, added in
quadrature.

###  2.2 Spectroscopic observation

Low-resolution optical spectroscopic observations of SN 2017iro were obtained
at 34 epochs during 2017 December 01 (JD 2458089.4) to 2018 July 05 (JD
2458305.2). Two grisms Gr#7 (3500–7800 Å) and Gr#8 (5200–9250 Å) available
with HFOSC, were used to obtain the SN 2017iro spectra. The spectra of arc
lamp and spectrophotometric standards were obtained along with the SN for
calibration purposes. The spectroscopic data reduction was carried out in a
standard way using IRAF. After correcting the observed frames for bias and
flat fields, one-dimensional spectra were extracted using the optimal
extraction method. Wavelength calibration was done by applying the dispersion
solutions obtained using arc lamp spectra. The night-sky emission lines were
utilized to secure the wavelength calibration, and small shifts were applied
wherever necessary. The spectrophotometric standard observed on the same night
was used in obtaining the instrumental response curves for flux calibration of
the spectra. For those nights where standard star observations were not
possible, the response curves obtained during nearby nights were used. To
construct a single flux calibrated spectrum, the spectrum in both the grisms
(Gr#7, Gr#8) were combined. The spectra were then scaled with respect to the
calibrated $UBVRI$ fluxes to bring them to an absolute flux scale. Finally,
the SN spectra were corrected for the host galaxy redshift of $z$ = 0.006191
(from NED).

Table 4: Estimated explosion epoch of SN 2017iro using different methods (see Section 3.1). Method adopted | Explosion epoch
---|---
Power law [$L(t)=A\,\times\,(t-t_{0})^{0.78}$] fit | JD 2458077.1 (2017 November 19.6)
Analytical fit (Taddia et al., 2018) | JD 2458080.0 (2017 November 22.5)
Toy model fit (Valenti et al., 2008) | JD 2458084.0 (2017 November 26.5)
Mean explosion epoch | JD 2458080.4 (2017 November 22.9)
Table 5: Light curve parameters of SN 2007iro (see Section 3.2). Filter | Max epoch | $\rm m_{max}$ | $\rm\Delta m_{15}$ | $\rm\Delta m_{40}$ | Late time slope | $\rm\Delta d_{0.25}$
---|---|---|---|---|---|---
| (JD) | (mag) | (mag) | (mag) |  ($\sim$40 – 140 d) |  ($>$140d) | (d)
|  |  |  |  |  (mag 100 d-1) |  (mag 100 d-1) |
$U$ |  2458095.7 $\pm$ 0.5 |  15.57 $\pm$ 0.01 |  2.41 $\pm$ 0.07 |  3.55 $\pm$ 0.01 | — | — |  7.70 $\pm$ 0.14
$B$ |  2458096.2 $\pm$ 0.5 |  15.94 $\pm$ 0.01 |  1.46 $\pm$ 0.05 |  2.44 $\pm$ 0.01 | 1.27 | — |  10.40 $\pm$ 0.14
$V$ |  2458098.0 $\pm$ 0.5 |  15.55 $\pm$ 0.01 |  0.77 $\pm$ 0.04 |  1.70 $\pm$ 0.01 | 1.73 | 1.80 |  11.90 $\pm$ 0.14
$R$ |  2458099.4 $\pm$ 0.5 |  15.14 $\pm$ 0.01 |  0.69 $\pm$ 0.02 |  1.54 $\pm$ 0.01 | 1.57 | 1.52 |  12.20 $\pm$ 0.14
$I$ |  2458100.5 $\pm$ 0.5 |  14.97 $\pm$ 0.01 |  0.52 $\pm$ 0.02 |  1.14 $\pm$ 0.01 | 1.75 | 1.90 |  14.70 $\pm$ 0.14
Bolometric |  2458095.7 $\pm$ 0.5 |  16.12 $\pm$ 0.01 |  1.05 $\pm$ 0.04 |  1.89 $\pm$ 0.01 | 1.65 | — |  10.40 $\pm$ 0.14

Figure 2: Left panel: The observed $UBVRI$ light curves of SN 2017iro. For
clarity, the light curves in different bands have been shifted vertically by
the indicated amount. Right panel: The $V$-band magnitude of SN 2017iro is
over-plotted with other Type Ib events i.e., iPTF13bvn, SN 2012au, SN 2009jf,
SN 2007Y, and SN 1999dn (see Table 6 for references). The light curves of
other events have been shifted arbitrarily to match the date of maximum and
the peak magnitude to SN 2017iro. The evolution of light curves near the peak
is also shown in the inset.

##  3 Photometric evolution

###  3.1 Estimation of explosion epoch and optical light curves

The last non-detection of SN 2017iro was reported (Wiggins, 2017) on 2017
November 15.1 (JD 2458072.6) with a limiting magnitude of 19 mag. However, the
SN was discovered on 2017 November 30.5 (JD 2458088.0). Since there is a large
gap between the first detection and last non-detection various approaches were
considered for estimating the explosion epoch. The explosion epoch was
estimated by fitting a power law of the form
$L(t)=A\,\times\,(t-t_{0})^{0.78}$ (Piro & Nakar, 2013) to the pre-maximum
phase of the LC. Here, $L$ denotes the luminosity at time $t$, $A$ is a
constant that defines rising rate, and $t_{0}$ is the time of the explosion.
The best fit value of $t_{0}$ is computed as JD 2458077.1 (2017 November
21.1). While evaluating the explosion parameters of SN 2017iro (see Section
3.4), the model of Valenti et al. (2008) was best fitted for an explosion date
of JD 2458084.0 (2017 November 26.5). Analytical function (Taddia et al.,
2018) fitted to the LC (see Section 3.2) yielded an explosion epoch of JD
2458080.0 (2017 November 22.5). An average of the above values, i.e., JD
2458080.4 $\pm$ 2.0 (2017 November 22.9), was adopted as the explosion epoch
of SN 2017iro (see Table 4). Further, the typical rise times of Type Ib SNe in
$R$-band and bolometric light curve are $\sim$ 22 d (Taddia et al., 2015) and
$\sim$17 d (Lyman et al., 2016), respectively. The $R$-band and the bolometric
peak of SN 2017iro occurred on JD 2458099.4 and JD 2458096.6, respectively
(see Section 3.2). The mean explosion date estimated above is consistent with
the rise time criteria.

The $UBVRI$ light curves of SN 2017iro are shown in Fig. 2 (left panel). A
well-sampled light curve covering the pre-maximum to nebular phase was
obtained using the photometric observation made on 40 epochs. In the right
panel of Fig. 2, the $V$-band light curve of well-studied Type Ib SNe:
iPTF13bvn (Srivastav et al., 2014; Kuncarayakti et al., 2015b), SN 2012au
(Pandey et al., 2021), SN 2009jf (Sahu et al., 2011; Valenti et al., 2011), SN
2007Y (Stritzinger et al., 2009), and SN 1999dn (Benetti et al., 2011), having
observations during pre-maximum and nebular phases are plotted along with SN
2017iro for comparison. The light curve evolution and estimation of related
parameters are discussed in the following subsections.

###  3.2 Light curve parameters

Various light curve parameters were obtained by fitting the light curves in
different bands with an analytical function proposed by Taddia et al. (2018,
see their section 3). This method is based on the empirical model suggested by
Vacca & Leibundgut (1996) for thermonuclear events whose LCs are similar to
that of SE-SNe. The epoch of maximum and the light curve parameters of SN
2017iro estimated from the LCs of different pass-bands are listed in Table 5
and are compared with other SE-SNe. A lag of $\sim$5 d between $U$ and
$I$-band maximum is inferred for SN 2017iro. This duration is shorter in
comparison to $\sim$8 d for the slow-declining SN 2009jf (Sahu et al., 2011)
and $\sim$9 d for the fast-declining iPTF13bvn (Srivastav et al., 2014). The
early appearance of maximum in bluer bands is a common behaviour of SE-SNe,
and SN 2017iro follows this trend.

The luminosity decline rate parameter, $\rm\Delta m_{15}$ (difference in
magnitude at peak and 15 d later), is handy in the context of Type Ia SNe,
where it is found to be correlated with the peak luminosity. A smaller
$\rm\Delta m_{15}$ value favours more luminous objects (Phillips, 1993). Large
data sets obtained from different surveys have provided an opportunity to
investigate if such a correlation exists for SE-SNe. No direct correlation was
found between the light curve decline rate parameter and luminosity by Drout
et al. (2011) and Taddia et al. (2015). However, the sample study of a larger
set of SE-SNe in Taddia et al. (2018) hinted towards the fact that luminous
events tend to have broader light curves, i.e., a smaller value of $\rm\Delta
m_{15}$. In the following paragraph, we compared various light curve
parameters of SN 2017iro with some well-studied Type Ib SNe.

The $\rm\Delta m_{15}$ values of SN 2017iro are 1.46, 0.77, 0.69 and 0.52 mag
in $B$, $V$, $R$ and $I$ bands, respectively. These are similar to the values
inferred for SN 2012au (Pandey et al., 2021, 1.60, 0.90, 0.70, and 0.43 mag,
respectively). On the other hand, iPTF13bvn shows comparatively rapid
evolution (1.82, 1.16, 1.22 and 0.90 mag, respectively, Srivastav et al.,
2014) whereas SN 2009jf shows a relatively slower evolution (Sahu et al.,
2011, 0.91, 0.50, 0.31, and 0.31 mag, respectively). The above can be inferred
from the comparison of normalised $V$-band light curves shown in right panel
of Fig. 2. The comparison of early phase light curve indicates that the
decline rate of SN 2017iro is intermediate between the fast declining SN
2007Y, iPTF13bvn and the slow declining SN 2009jf. The intermediate width of
light curves of SN 2017iro could possibly indicate that ejecta mass of SN
2017iro lies in between the inferred ejecta mass of iPTF13bvn and SN 2009jf
(see Section 5.3).

The decay in the magnitude from the peak to +40 d is denoted by the $\rm\Delta
m_{40}$ parameter by Taddia et al. (2018). $\rm\Delta m_{40}$ was found to be
in the range of 1.5 to 1.7 mag (in $r$-band) for their sample of SE-SNe. SN
2017iro lies towards the lower end (1.54 mag in $R$-band) of the observed
range, indicating a relatively faster evolution for a SE-SN. The late phase
($>$ 40 d) light curve decay rate of SN 2017iro is also computed in all bands
(see Table 5). The decay rates between $\sim$ +40 to $\sim$ +140 d are 1.27,
1.73, 1.57 and 1.75 mag 100 d-1 (in $B$, $V$, $R$ and $I$ bands, respectively)
and beyond $\sim$140 d is 1.80, 1.52 and 1.90 mag 100 d-1 (in $V$, $R$ and $I$
bands, respectively). It is to be noted that usually, SE-SNe display faster
decline than Type II events during the late phases. Such features are
indicative of higher $\gamma$-rays escape. The implications of the fast
declining nature of SN 2017iro is discussed in Section 5.1.

The $V$-band absolute magnitude ($M_{V}$) of SE-SNe is known to display a wide
range between $\sim$ –16.5 to $\sim$ –19.5 mag (Richardson et al., 2006; Drout
et al., 2011; Taddia et al., 2018). To estimate the $M_{V}$ of SN 2017iro, we
adopted distance of the host galaxy NGC 5480 as 30.8 $\pm$ 2.2 Mpc (corrected
for Virgo infall only) and the corresponding distance modulus 32.44 $\pm$ 0.15
mag (for an H0 = 73 km s-1 Mpc-1). The adopted distance is consistent with the
distance estimate of 30.73 Mpc using Tully-Fisher relation (Theureau et al.,
2007). After applying a correction for the reddening $E(B-V)$ = 0.28 mag (see
Section 3.3), the $M_{V}$ of SN 2017iro is derived as $-17.76\pm 0.15$ mag.
This indicates that SN 2017iro is fainter than SN 2007uy, SN 2009jf and SN
2012au but brighter than SN 2007Y and iPTF13bvn.

Figure 3: Colour evolution of SN 2017iro, plotted along with iPTF13bvn, SN
2012au, SN 2009jf, SN 2007uy and SN 2007Y. Both the MW and the host galaxy
extinction corrections have been applied for each SN. The $B-V$ colour
templates from Stritzinger et al. (2018) are over-plotted in the top panel for
comparison. The bibliographic sources are the same as mentioned in the text
(Section 3). Figure 4: The psuedo-bolometric light curve of SN 2017iro
(connected with a bold line) compared with other similar Type Ib events
(iPTF13bvn, SN 2012au, SN 2009jf, SN 2007uy, SN 2007Y and SN 1999dn). The
psuedo-bolometric light curves were obtained using SUPERBOL.

###  3.3 Reddening and colour evolution

The estimate of reddening suffered by SN is crucial as it affects the
explosion parameters which are derived using the observables. The reddening
due to ISM in the Milky Way along the line of sight (LOS) of the SN can be
inferred using the all-sky dust extinction map (e.g. Schlafly & Finkbeiner,
2011; Schlafly et al., 2016). The determination of reddening within the host
galaxy is a challenging task. In general, the sites of SE-SNe, are associated
with star-forming regions (e.g. Anderson et al., 2012; Eldridge et al., 2013)
and may suffer a considerable amount of reddening. Several photometric and
spectroscopic techniques have been proposed to determine the reddening within
the host galaxy but with caveats.

The Galactic reddening along the LOS of SN 2017iro as derived from the all-sky
dust extinction map (Schlafly & Finkbeiner, 2011) is $E(B-V)$ = 0.016 $\pm$
0.001 mag. The equivalent width (EW) of Na i D absorption lines is considered
as a good tracer of reddening within the host-galaxy (Barbon et al., 1990;
Turatto et al., 2003; Poznanski et al., 2012) but with caveats (Poznanski et
al., 2011; Phillips et al., 2013). To estimate the EW of Na i D, we stacked
the low-resolution spectra of SN 2017iro obtained near the maximum light. A
reasonably strong Na i D absorption feature is seen at the redshift of the
host galaxy with an EW = 1.34 $\pm$ 0.10 Å. An average value of
$E(B-V)_{host}$, computed from the relation of Barbon et al. (1990) and
Turatto et al. (2003) is 0.27 $\pm$ 0.04 mag. For further analysis we adopt a
total extinction $E(B-V)_{MW+host}$ = 0.28 $\pm$ 0.04 mag (Milky Way + host)
and Cardelli et al. (1989) extinction law with ratio of selective to total
extinction $R_{V}$ =3.1.

In Fig. 3, the $B-V$, $V-R$, and $R-I$ colours of SN 2017iro are compared with
other well-studied Type Ib events (see Table 6 for details and references). It
is seen that the overall shape of colour evolution for each colour is similar
(also see Stritzinger et al., 2018). Up to $\sim$ +20 d, the $B-V$ colour
evolves faster than the $V-R$ and $R-I$ colours. During the late phase ($>$ 20
d), the colour evolution is almost flat.

Table 6: Sample of Type Ib SNe used in this study.

Name | $E(B-V)_{tot}$ | Distance | References
---|---|---|---
| (mag) | (Mpc) |
SN 1999dn | 0.10 | 38.90 | 1
SN 2007Y | 0.11 | 19.31 | 2
SN 2007uy | 0.63 | 29.50 | 3
SN 2009jf | 0.11 | 34.25 | 4, 5
SN 2012au | 0.06 | 23.50 | 6, 7
iPTF13bvn | 0.22 | 22.49 | 8, 9

References: 1 (Benetti et al., 2011), 2 (Stritzinger et al., 2009), 3 (Roy et
al., 2013), 4 (Sahu et al., 2011), 5 (Valenti et al., 2011), 6 (Milisavljevic
et al., 2013a), 7 (Takaki et al., 2013), 8 (Srivastav et al., 2014), 9
(Kuncarayakti et al., 2015b).

Figure 5: The blackbody temperature and radius evolution of SN 2017iro. These
parameters were derived from the photometric data using SUPERBOL. The shaded
regions indicate 1-$\sigma$ uncertainties.

Figure 6: Pre-maximum along with early post-maximum (–7 to +9 d) and late
post-maximum (+18 to +53 d) phases spectral evolution of SN 2017iro are shown
in the top and bottom panels, respectively. Prominent spectral lines are
marked. The labelled phases are with respect to the $B$-band maximum (JD
2458096.2). The spectra have been corrected for redshift and reddening.
Features arising from the host are marked with dashed vertical lines. Major
telluric bands are shown with a circled plus symbol. The spectra shown in
Figures 6, 8 and 9 are available as data behind the figure. Figure 7: SN
2017iro spectral evolution compared with other Type Ib events (iPTF13bvn, SN
2012au, SN 2009jf, SN 2008D and SN 2007uy) at similar epochs. Panels A, B and
C represent the pre-maximum, near-maximum and post-maximum spectra,
respectively. Telluric lines are shown with circled plus symbols.

###  3.4 Quasi-bolometric light curve

The quasi-bolometric light curve of SN 2017iro was constructed using the
latest version of Python-based code SuperBol (Nicholl, 2018). The $UBVRI$
magnitudes after correcting for the total extinction $E(B-V)$ = 0.28 mag, and
adopted distance to NGC 5480 were given as input parameters. The $U$ and $B$
band magnitudes were extrapolated by assuming a constant colour during the
late phases. The flux integration was performed only over the optical
wavelengths. The resultant quasi-bolometric light curve of SN 2017iro along
with well studied Type Ib events (estimated similarly) are shown in Fig. 4.
Furthermore, the black body fit parameters (temperature and radius) which is a
default outcome of the code is also plotted in Fig. 5.

We fitted Taddia et al. (2018) equation to the quasi-bolometric light curve of
SN 2017iro and estimated parameters listed in Table 5. The light curve peaked
at JD 2458095.7 with log${}_{10}L$ = 42.39 $\pm$ 0.09 erg s-1. The $\rm\Delta
m_{15_{bol}}$ value of SN 2017iro is 1.05 mag, which is higher than that of SN
2009jf (0.60 mag) but is comparable to SN 2007Y (0.80 mag). The decay rates
after +40 d of SN 2017iro (0.016 mag d-1) and SN 2009jf (0.014 mag d-1) are
similar implying relatively faster evolution than the cobalt decay rate
($\sim$0.01 mag d-1, see also Section 5.1). This decay rate is consistent with
the range 0.014 – 0.018 mag d-1 seen in Type Ib events (Taddia et al., 2018).

Figure 8: Late post-maximum and early nebular phase ($\sim$+67 to +112 d)
spectral evolution of SN 2017iro. Prominent spectral lines are also marked.
Other descriptions are similar as mentioned in Fig. 6. The spectra shown in
Figures 6, 8 and 9 are available as data behind the figure. Figure 9: Nebular
phase ($\sim$+134 to +209 d) spectra of SN 2017iro. Various nebular emission
lines are shown. The dashed vertical lines indicate the features arising from
the host and the telluric bands are shown with a circled plus symbol (for
other descriptions, see Fig. 6). The spectra shown in Figures 6, 8 and 9 are
available as data behind the figure. Figure 10: Comparison of nebular phase
spectral features of SN 2017iro with other well studied Type Ib events. The SN
references are same as described in Section 4.1 in addition to Master OT
J120451.50+265946.6 (MOT, Singh et al., 2019).

Figure 11: Evolution of the velocity of different spectral lines observed in
SN 2017iro. Error in these estimates was never below 800 km s-1. The line
velocities of different SNe (SN 2007Y, 2008D, 2009jf, 2012au and iPTF13bvn)
are also compared.

##  4 Spectroscopic evolution

The early (photospheric phase) spectra are useful to examine the outermost
regions of the expanding ejecta. For example, the presence/absence of residual
hydrogen and/or helium during the pre-explosion stellar evolution can be
explored with the early phase spectra. In the nebular phase, because of
expansion, the ejecta becomes optically thin, and hence spectra during this
phase provide essential information about the inner layers of the ejecta.

In this section, the spectral evolution of SN 2017iro covering pre-maximum to
nebular phase is presented. The densely sampled spectroscopic data provided an
opportunity to compare important spectral features with other well-studied
events. All spectra were corrected for the heliocentric velocity of 1856 km
s-1 of the host galaxy (NGC 5480). The evolution of spectroscopic features is
discussed in the following sections.

###  4.1 Pre-maximum and early post-maximum phase

The spectral evolution in the pre-maximum and around maximum phase is shown in
Fig. 6 (top panel). During early phase (i.e. pre-maximum), spectroscopy of SN
2017iro was done almost every night. The first spectrum at $\sim$–7 d clearly
shows the broad absorption feature of He i 5876 Å line, and weak signatures of
other He i lines. He i 6678 Å is blended with narrow H${\alpha}$ line from the
host galaxy, whereas the He i 7065 Å feature lies at the edge of the telluric
absorption band at 6870 Å. Other narrow lines (H${\beta}$ 4861 Å, [N ii] 6584
Å and [S ii] 6717, 6731 Å) originating from the underlying H ii region are
also clearly visible. In the extreme blue region of the observed spectra, well
developed Ca ii H & K lines (3934 and 3968 Å) are seen. The spectrum also
shows features of Fe ii triplet (4924, 5018 and 5169 Å), Si ii 6355 Å and a
weak absorption due to Ca ii NIR triplet (8498, 8542 and 8662 Å).

In the pre-maximum spectral evolution, the region between 4000 to 4300 Å
appears flat, possibly due to the blending of metal lines in the spectra. The
absorption around 4400 Å is possibly a blend of He i, Fe ii and Mg ii. As the
SN evolves towards the maximum, the lines due to He i, Fe ii, and Si ii become
more prominent. In the early post-maximum spectra, the Si ii 6355 Å feature
starts weakening and vanishes in the spectra obtained on +18 d (cf. Fig. 6,
bottom panel). A feature around 5500 Å begins to develop in the spectrum of
$\sim$+6 d, which is likely a blend of Fe ii 5535 Å and Sc ii 5527 Å lines.
The Na i D lines from the host galaxy are seen as a narrow absorption in the
spectra.

The spectra of SN 2017iro during the pre-maximum (Fig. 7, Panel: A) and close
to maximum phase (Fig. 7, Panel: B) are compared with the spectra of few other
well-studied Type Ib SNe SN 2009jf (Sahu et al., 2011), iPTF13bvn (Srivastav
et al., 2014; Kuncarayakti et al., 2015b), SN 2007Y (Stritzinger et al.,
2009), SN 2008D (Modjaz et al., 2009b) and SN 2012au (Pandey et al., 2021)
around similar epochs. A good similarity can be noticed in the spectrum of SN
2017iro obtained at –7 d and near maximum with those of SN 2009jf and
iPTF13bvn. Except for iPTF13bvn, the Fe ii line near 5000 Å is found to be
prominent in other SNe used in comparison (cf. panel: A in Fig. 7). However,
in SN 2017iro, the Fe ii line appears to be contaminated by narrow H${\beta}$
line from the underlying H ii region. The spectrum of SN 2012au is shifted
blue-wards (cf. panel: B in Fig. 7) due to its higher expansion velocity
(Milisavljevic et al., 2013a; Takaki et al., 2013).

###  4.2 Late post-maximum and early nebular phase

The spectral evolution during +18 to +53 d is displayed in the bottom panel of
Fig. 6. The continuum becomes redder during this phase, and the Ca ii H & K
lines disappear, and other metallic lines appear in this part of the spectrum.
However, the region below 4500 Å in our spectrum, is noisy, so lines could not
be identified. The He i lines are initially prominent, but beyond +45 d, they
weaken. In the blue wing of He i 5876 Å line, lines due to Sc ii start
emerging. The Fe ii lines around 5000 Å also start blending. The Ca ii NIR
triplet gradually transformed into emission dominated. Beyond +40 d, when the
light curve enters into the exponential decline, the forbidden lines [O i]
6300, 6364 Å and [Ca ii] 7291, 7324 Å start appearing.

From the light curve, it is clear that the supernova entered into the linear
decline phase beyond 50 d after the explosion, which indicates a transition
from the photospheric phase to the nebular phase. The emergence of nebular
lines, especially the forbidden lines of [O i] 6300, 6364 Å and [Ca ii] 7291,
7324 Å also signifies the transition to the nebular phase. The spectrum taken
during +67 to +112 d is displayed in Fig. 8. During this phase, the spectra
are slowly getting dominated by the forbidden emission lines. The nebular line
[O i] 5577 Å is also seen along with O i 7774 Å line, which is probably
powered by an oxygen recombination cascade. The Ca ii NIR triplet is dominated
by the emission component and possibly blended with the [C i] 8730 Å line. The
spectrum obtained on +112 d still shows the presence of continuum and is not
fully nebular. In Fig. 7 (panel: C), the spectra near a month after the
maximum light is compared. The $\sim$+31 d spectrum of SN 2017iro is similar
to the spectrum of iPTF13bvn and SN 2008D. Also, He i 6678, 7065 Å lines in SN
2009jf are weaker as compared to other SNe. It is also worth mentioning that
between 4000 –5000 Å SN 2017iro and SN 2008D exhibit comparatively flat
structures however, iPTF13bvn and SN2009jf have strong features in this
wavelength range.

###  4.3 Nebular phase

The spectra taken during +134 to +209 d is presented in Fig. 9. The spectra
are dominated by the forbidden emission lines of [O i] and [Ca ii], narrow
lines from the nearby H ii region are also clearly seen. The [O i] 5577 Å
which was fairly strong in the early nebular phase (before $\sim$ 100 d),
faded after that and almost vanished in the nebular spectra. The semi-
forbidden line Mg i] 4571 Å is found to be always weak in the spectra of SN
2017iro. The emission component of the Ca ii NIR triplet, which was stronger
till $\sim$ 120 d starts weakening, and by $\sim$ 176 d it merges with the
almost flat continuum. The forbidden [O i] and [Ca ii] lines show complex
profile.

In Fig. 10, the nebular phase spectrum of SN 2017iro at +209 d is compared
with those of other type Ib events, obtained at similar epoch. The emission
features of semi-forbidden Mg I] 4571 Å, forbidden [O I] 6300, 6364 Å, and [Ca
II] 7291, 7324 Å are prominently visible in these SNe. It is interesting to
note that in SN 2017iro, the strength of [O I] and [Ca II] lines are similar
whereas, in SN 2009jf, SN 2012au and Master OT J120451.50+265946.6 (MOT) the
[O I] line is stronger than the [Ca II] line. In SN 2017iro, the Mg I] is
comparatively very weak than other SNe. The ratio of Mg I] to [O I] lines in
SN 2017iro is estimated as $\sim$0.15 which is similar to MOT ($\sim$0.3,
Singh et al., 2019) but lesser than iPTF13bvn ($\sim$0.85, Kuncarayakti et
al., 2015b).

Amongst the nebular forbidden emission lines, the Mg I] line lies in the bluer
and crowded region of the spectra, where generally signal is poor.
Furthermore, Ca is a trace element and its distribution is not necessarily
representative of the bulk of the ejecta. Hence, the [O I] doublet which
emerges in relatively clear region, is generally used as the tracer of the
explosion geometry (Mazzali et al., 2005; Maeda et al., 2008a; Modjaz et al.,
2008a; Taubenberger et al., 2009; Milisavljevic et al., 2010). From the Fig.
10, similarity between nebular spectrum of SN 2009jf and SN 2017iro is evident
with [O I] line profile being asymmetric and multi-peaked in both the objects.
However, asymmetry in [O I] line profile of SN 2017iro is weaker as compared
to SN 2009jf with the peaked component of [O I] line centered at rest
wavelength. The asymmetric and multi-component [O I] profile, seen in the
nebular spectra, can be reproduced with a complex ejecta geometry of an
aspherical explosion (Mazzali et al., 2005; Maeda et al., 2007a). The observed
features in the asymmetric [O I] line profile in SN 2017iro, could be
explained with a high density core of the ejecta (Taubenberger et al., 2009;
Maeda et al., 2008a).

The gas-phase oxygen abundance in the vicinity of supernova location is
estimated using the relation of Pettini & Pagel (2004). N2 and the O3N2
indices were computed from the last spectrum obtained on $\sim$ 209 d. The
average value of oxygen abundance 12 + log(O/H) (estimated using O3N2 and N2
indices) for the supernova region is found to be 8.64 $\pm$ 0.2 which is
remarkably similar to iPTF13bvn (8.63, Kuncarayakti et al., 2015b). In a
recent study under the CALIFA (Sánchez et al., 2012) survey program, Galbany
et al. (2016) used wide-field integral field spectroscopy to compute the
global metallicity of NGC 5480 and metallicity at the location of SN 1988L.
The estimated values are 8.58 and 8.55 for the local and global, respectively.
This shows that our estimate of the oxygen abundance is in fairly good
agreement with the previous study, and is close to the solar value of 12 + log
(O/H) = 8.69 $\pm$ 0.05 (Asplund et al., 2009). Further, the computed
metallicity close to the SN location is similar to the estimates of Modjaz et
al. (2011, 8.49 $\pm$ 0.19) and Sanders et al. (2012, 8.48 $\pm$ 0.16), made
at the location of other SNe Ib.

###  4.4 Line velocities

The velocities of lines Fe ii 5169 Å, He i 5876 Å, He i 6678 Å and Si ii 6355
Å were estimated by fitting Gaussian profile to the respective absorption
trough after correcting the spectra for the redshift of host galaxy of SN
2017iro. The estimated line velocities of SN 2017iro along with SNe 2007Y,
2008D, 2009jf, 2012au, and iPTF13bvn collected from the literature are plotted
in Fig. 11. In case of SN 2017iro, the He i 5876 Å velocity is $\sim$11500 km
s-1 in the beginning and declined to $\sim$9000 km s-1 near maximum. In the
post-maximum phase, it evolves slowly and becomes flat ($\sim$7000 km s-1)
after +70 d. All the SNe (except SN 2007Y), used in the comparison, show a
steep decline in the He i 5876 Å line velocity during the pre-maximum phase.
The post-maximum evolution of He i shows a modest increase of $\sim$900 $\rm
km\ s^{-1}$ giving rise to a bump like feature in the velocity evolution. A
similar feature can also be seen in SN 2009jf, iPTF13bvn, and SN 2008D during
maximum to $\sim$+20 d. The origin of this feature needs to be explored. The
evolution of He i 5876 Å line velocity beyond $\sim$+25 d, is similar in SN
2017iro, iPTF13bvn, and SN 2009jf but lower than SN 2012au. In homologous
expanding SNe ejecta, it is difficult to accurately determine the photospheric
velocity from the observed spectral features as none measurable feature of the
spectra is connected directly to photospheric velocity (Takáts & Vinkó,
2012).However, the prominent absorption features produced due to Fe ii 5169 Å
can be used as a good tracer of SN photospheric velocity as it is a less
optically thick line compared to several others lines (Dessart & Hillier,
2005). The Fe ii line velocity of SN 2017iro is compared in the middle panel
of Fig. 11. The ejecta velocity measured using Fe ii line of all the SNe shows
a steep decline till $\sim$+25 d and flattens thereafter. The Fe ii line
velocity of SN 2017iro near-maximum is comparable to SN 2007Y, SN 2009jf, and
iPTF13bvn, but slower than SN 2008D and SN 2012au. In the later phase ($>$+25
d), the Fe ii line velocity of SN 2017iro is in between SN 2009jf and SN
2007Y. The Fe ii velocity around maximum light ($\sim$9000 $\pm$ 800 km s-1)
is used to estimate the explosion energy of SN 2017iro (see Section 5.3).

##  5 Discussion

The observed fraction of Type Ib SNe is low, they are rare objects (Li et al.,
2011b; Graur et al., 2017a, b; Shivvers et al., 2019). There are only a
handful events for which studies covering both photospheric and nebular phases
are available. A detailed investigation of SN 2017iro, hence is a good
addition to this sample as it fills the gap between fast declining (e.g. SN
2007Y, iPTF13bvn) and slow declining (e.g. SN 2009jf) objects (see Section
3.2). Here, we discuss the nature of SN 2017iro based on some significant
light curve and spectral features, constrain its progenitor mass, and also
estimate the explosion parameters.

###  5.1 Late phase light curve heterogeneity

The light curve of SE-SNe follows a linear decline with the onset of the
nebular phase (usually $>$ 100 days after the explosion). This is a
consequence of energy injection from the radioactive decay of 56Co
$\rightarrow$ 56Fe. The efficiency of $\gamma$-ray trapping within the ejecta
affects the decay rate. In general, it is found that the late phase light
curve of SE-SNe declines faster than the expected decay rate of 0.98 mag (100
d)-1 from 56Co $\rightarrow$ 56Fe transition. This is primarily due to the
lower opacity of the ejecta for the $\gamma$-rays (i.e., incomplete trapping).
Therefore, the light curve decay rate during the nebular phase can be used to
infer the ratio of explosion energy, and mass of the ejecta. The late phase
(beyond 140 d) decline rate of SN 2017iro is estimated in different bands and
are listed in Table 5. The light curves are found to decline significantly
faster than 0.98 mag (100 d)-1.

An investigation of the late time light curve of SE-SNe by Wheeler, Johnson, &
Clocchiatti (2015) revealed the diverse nature of their evolution, indicating
dispersion in the ejecta masses and kinetic energies. Our comparison of the
late-time decay rate of well studied Type Ib SNe (cf. Fig. 2: right panel)
indicates that there exist two groups of events. SN 1999dn and SN 2012au fall
under the first category in which the late phase decline rate is significantly
slower than the 56Co $\rightarrow$ 56Fe decay rate. The slower light curve
decline rate indicates the need for an additional source of energy which could
be interaction of the ejecta with the circumstellar material (CSM), a magnetar
or presence of light echo. Wheeler et al. (2015) suggested the presence of CSM
in SN 1999dn. The presence of smoothly distributed CSM around SN 2012au was
inferred from its detection in radio wavebands (Kamble et al., 2014). However,
very late nebular spectrum of SN 2012au, obtained $\sim$ 6.2 years after the
explosion, indicated magnetar/pulsar wind nebulae as the most probable energy
source (Milisavljevic et al., 2018). In the second group of SE-SNe, the light
curve decline rate is faster than the expected decline rate of 56Co – 56Fe.
The faster decline rate could be due to incomplete trapping of $\gamma$-rays
produced via 56Co – 56Fe decay, in the SN ejecta. The late-time decay rate of
SN 2017iro indicates it falls in this category and is similar to that of SN
2009jf, SN 2007Y and iPTF13bvn. It is worth mentioning that during the late
nebular phase, an increasing fraction of photons emitted at IR wavelength may
also steepen the decline rate in optical bands.

###  5.2 56Ni mixing

The lines of He i, the identifying features of Type Ib SNe, are seen at a
velocity of $\sim$ 11500 km s-1 in the first spectrum ($\sim$–7 d) of SN
2017iro (cf. Fig. 11). The origin of these lines require non-thermal
excitation and ionization (Lucy, 1991; Swartz, 1991). The accelerated
electrons responsible for the non-thermal excitation are regulated by the
$\gamma$-rays produced in the radioactive decay of newly synthesized 56Ni
(Harkness et al., 1987). It is shown that the mere presence of an appreciable
amount of helium in the progenitor would not lead to the presence of helium
lines in the spectra (Dessart et al., 2012). To excite He i at high
velocities, $\gamma$-rays need to be in the proximity of the helium layer,
either due to the leakage of $\gamma$-rays through the inner ejecta or by
substantial mixing of 56Ni up to the outer ejecta.

Mixing may arise due to various reasons such as large-scale asymmetric
explosion jets and asymmetry of the propagating shock etc. (see, Maund et al.,
2009; Couch et al., 2011; Dessart et al., 2012, and references therein). The
effect of 56Ni mixing on the light curves and spectra of SE-SNe have been
investigated in detail by several authors. Theoretical studies (Shigeyama et
al., 1990; Hachisu et al., 1991) argued that 56Ni mixing within the helium
layer might depend on the progenitor composition. The degree of 56Ni mixing
may also significantly influence the overall light curve shape (including the
post-maximum phase) and the efficient formation of He i spectral lines (Ensman
& Woosley, 1988; Woosley & Eastman, 1997; Shigeyama et al., 1990; Cano et al.,
2014; Dessart et al., 2015, 2016; Taddia et al., 2018).

The early time colour evolution of Type Ib/c SNe can also be used as a
diagnostic for the 56Ni distribution in the ejecta. It is shown that in the
case of strong mixing, the colour curves of Type Ib/c SNe monotonically redden
during the photospheric phase, because of continuous and progressive effects
of 56Ni heating. Whereas, in the case of weak to moderate mixing, the colour
evolution does not show monotonic behaviour due to delayed heating from 56Ni
(Yoon et al., 2019, and references therein). The early phase $B-V$ colours of
different Ib events along with 56Ni mixing models (Yoon et al., 2019) are
plotted in Fig. 12. The colour evolution of SN 2007Y, SN 2009jf, and iPTF13bvn
follow a non-monotonic trend which is interpreted in the framework of weak to
moderate mixing of 56Ni in the ejecta. However, the non-availability of
sufficiently early phase data points for SN 2017iro prevents us from arriving
at any firm conclusion on the role of mixing.

###  5.3 Physical parameters of explosion

The peak of the light curve and its shape in SE-SNe are mainly regulated by
the synthesized radioactive 56Ni, kinetic energy ($E_{\rm k}$) and ejecta mass
($M_{\rm ej}$). Estimation of these parameters are crucial in understanding
the explosion properties. The explosion parameters can be derived employing
either detailed hydrodynamical modelling of both the light-curve and the
spectra or using semi-analytical models. In the present study, several semi-
analytical formulations were used to fit the quasi-bolometric light curve (see
Section 3.4) of SN 2017iro.

Figure 12: The $B-V$ colour evolution of SN 2017iro compared with the
synthetic colour curves of SN Ib models by Yoon et al. (2019). The five bold
curves indicate varying degrees of 56Ni mixing (fm) as shown in the legend
(Y19). An increasing value of fm implying a stronger 56Ni mixing (for more
details, see Yoon et al., 2019).

####  5.3.1 Estimation of 56Ni mass, ejecta mass and kinetic energy

The optically thick phase of the light curve (photospheric phase), was fitted
by the formulation originally proposed by Arnett (1982) and later updated by
Valenti et al. (2008). Major assumptions in this analytical approach comprise
a small radius of the progenitor, spherical symmetry and homologous expansion
of the ejecta, constant opacity ($\kappa_{\rm opt}$) and a centrally located
and unmixed 56Ni (see Arnett, 1982; Valenti et al., 2008; Cano, 2013). The
parameters $M_{\rm Ni}$ (nickel mass) and $\tau_{\rm m}$ (diffusion time-
scale) were kept as free variables while obtaining the fit. Considering a
uniform density medium, the ejecta kinetic energy $E_{\rm k}$ and $\tau_{\rm
m}$ are, given by

$\tau_{\rm m}=\sqrt{2}\left(\frac{\kappa_{\rm opt}}{\beta
c}\right)^{1/2}\left(\frac{M_{\rm ej}}{v_{\rm ph}}\right)^{1/2},$ (1)

$E_{\rm k}\approx\frac{3}{5}\frac{M_{\rm ej}v^{2}_{\rm ph}}{2},$ (2)

where $\beta\approx 13.8$ is a constant of integration (Arnett, 1982) and c is
the speed of light. The optical opacity $\kappa_{\rm opt}$ is adopted as 0.07
cm2 g-1 (e.g. Chugai, 2000; Toy et al., 2016; Taddia et al., 2018). The quasi-
bolometric light-curve of SN 2017iro until 30 d from the explosion epoch was
fitted using least-square optimization. The best fit was obtained for a
$M_{\rm Ni}$ = 0.09 $\pm$ 0.04 M⊙ and $\tau_{\rm m}$ = 9.95 $\pm$ 0.44 d. The
derived 56Ni mass was further estimated using the bolometric rise time of 16.5
d and a peak bolometric luminosity log${}_{10}L$ = 42.39 erg s-1. The relation
given by Stritzinger & Leibundgut (2005) gives a value of 0.10 $\pm$ 0.01 M⊙.
Using the measured Fe ii line velocity 9000 km s-1 near the bolometric maximum
(see Section 4.1), the $M_{\rm ej}$ and $E_{\rm k}$ are estimated as 1.39 M⊙
and 0.75 $\times$ 1051 erg, respectively.

While investigating the physical parameters of SN 2002ap, Vinkó et al. (2004)
proposed analytic models for post-maximum phase with assumptions as described
in the beginning of this sub-section. In brief, these models are based on the
ejecta density configuration characterized by $x_{0}$ (fractional radius of
the core) and $n$ (power-law exponent) parameters. We used their model C1
which represents a ‘core-shell’ structure with a fixed density core of
fractional radius, $x_{0}$ = 0.15 and a fixed value of $\kappa_{\gamma}$ =
0.027 cm2 g-1 for grey atmospheres (Sutherland & Wheeler, 1984). The best-fit
model is shown in Fig. 13 with an $M_{\rm ej}$ = 4.3 M⊙, $M_{\rm Ni}$ = 0.05
M⊙, $E_{\rm k}$ = 1.85 foe and $v_{\rm max}$ = 8,500 $\rm kms^{-1}$.

Energy production rates of 56Ni $\rightarrow$ 56Co $\rightarrow$ 56Fe decay
chain computed by Nadyozhin (1994) were further used to examine the mass of
synthesized 56Ni during the explosion. The energy production curves for three
different values (i.e. 0.06, 0.10, and 0.15 M⊙) of the mass of 56Ni were over-
plotted in Fig. 13. The maximum in the quasi-bolometric light curve describes
roughly the transition between the emission deficit due to the large optical
depth (larger diffusion timescale) and an excess emission due to the stored
radiation post-maximum. The early post-maximum decline can be approximated by
the instantaneous energy production rate of 56Ni $\rightarrow$ 56Co decay.
Hence, we infer that a 0.10 M⊙ of $\rm{}^{56}Ni$ was synthesised in the
explosion due to its best-match to the initial post-maximum decline of the
quasi-bolometric light-curve of SN 2017iro. The large discrepancy between the
energy production rate and the quasi-bolometric luminosity during the late
phase is arising due to the escape of gamma-rays.

Figure 13: The quasi-bolometric light curve of SN 2017iro fitted with
analytical model C1 of Vinkó et al. (2004). Three curves (dotted, dashed, and
dash-dotted) representing the rate of energy production for different masses
of 56Ni synthesized during the explosion, based on the analytical formulation
by Nadyozhin (1994) have also been shown.

Based upon a study of 38 SE-SNe, Lyman et al. (2016) proposed a fitting
relation to estimate $M_{\rm Ni}$ (see their Eq. 4). To populate this data
sample, we collected 9 additional SE-SNe (5 IIb, 2 BL-Ic, 1 Ic-transitional,
and 1 Ib events) with well-sampled light curves from the literature (see Table
7). Also, to estimate Mbol, the data available only in the optical wavelength
was used. In Fig. 14, a correlation between Mbol and 56Ni mass is shown for
all the 47 events. The plot also indicates that among the population of SE-SNe
that the Type IIb events produce a lesser amount of 56Ni. However, there
appears to be no significant difference in 56Ni production between Type Ib and
Ic SNe, whereas Type Ic-BL SNe, in general, are brighter with higher 56Ni
mass. This is in line with other recent studies (Anderson, 2019; Prentice et
al., 2019). With a $M_{bol_{peak}}$ of –17.19 $\pm$ 0.22 mag (cf. Table 5) and
56Ni mass of 0.08 M⊙ (average value estimated from above methods), SN 2017iro
(shown with a star symbol in the Fig. 14) lies near the expected correlation
for SE-SNe. It is to be noted here that the explosion parameters inferred
using empirical relations are approximate, and detailed hydrodynamical
modelling (e.g. Taddia et al., 2018) would be required for better estimates.

Table 7: In addition to Lyman et al. (2016) following events are used to
estimate correlation between Mbol and 56Ni mass (see Section 5.3).

Name | Type |  Mbol |  56Ni mass | References
---|---|---|---|---
|  | (mag) |  (M⊙) |
SN 2011ei | IIb | – 15.8 |  0.03 $\pm$ 0.01 | 1
SN 2011fu | IIb | – 17.5 |  0.21 $\pm$ 0.02 | 2
SN 2012ap | Ic-BL | – 17.2 |  0.12 $\pm$ 0.02 | 3
SN 2012au | Ib | – 17.8 |  0.26 $\pm$ 0.04 | 4, 5, 6
SN 2013df | IIb | – 16.2 |  0.10 $\pm$ 0.03 | 7, 8
SN 2014ad | Ic-BL | – 18.1 |  0.27 $\pm$ 0.06 | 9
SN 2015as | IIb | – 16.7 |  0.08 $\pm$ 0.01 | 10
SN 2016coi |  Ic/Ic-BL∗ | – 16.9 |  0.11 $\pm$ 0.01 | 11
SN 2016gkg | IIb | – 17.4 |  0.09 $\pm$ 0.01 | 12

References: 1 (Milisavljevic et al., 2013b), 2 (Kumar et al., 2013), 3
(Milisavljevic et al., 2015), 4 (Milisavljevic et al., 2013a), 5 (Pandey et
al., 2021), 6 (Takaki et al., 2013), 7 (Morales-Garoffolo et al., 2014), 8
(Szalai et al., 2016), 9 (Sahu et al., 2018), 10 (Gangopadhyay et al., 2018),
11 (Kumar et al., 2018), 12 (Tartaglia et al., 2017b).
∗ Ic/Ic-BL transitional event.

Figure 14: Correlation between Mbol and 56Ni mass of 47 SE-SNe (including the
sample from Lyman et al. (2016) and 9 additional events, see Section 5.3).
Type Ib, Ic, Ic-BL and IIb are shown with different symbols. The bold
continuous line is best fit. The dark–pink and light–pink shaded regions are
indicative of 1$\sigma$ and 3$\sigma$ confidence intervals of the fit,
respectively. Table 8: Oxygen mass of Type Ib SNe.

Name |  Oxygen mass (M⊙) | References
---|---|---
SN 1996N | 0.11 – 0.21 | 1
SN 2007Y | 0.20 | 2
SN 2009jf | 1.34 | 3
iPTF13bvn | 0.33 | 4
MOT∗ | 0.90 | 5
SN 2017iro | 0.35 | This study

∗ Master OT J120451.50+265946.6
References: 1 (Sollerman et al., 1998), 2 (Stritzinger et al., 2009), 3 (Sahu
et al., 2011), 4 (Kuncarayakti et al., 2015b), 5 (Singh et al., 2019)

####  5.3.2 Oxygen mass

The [O I] line flux in the nebular spectra of SE-SNe can be used to estimate
the neutral oxygen mass that is required to produce the emission. In the high
electron density ($n_{e}\,\geq 10^{-6}$ cm-3) region, the minimum O mass can
be estimated by the relation $M_{O}=10^{8}\times D^{2}\times
F([\rm{OI}])\times exp{(2.28/T_{4})}$, provided by Uomoto (1986). Here, mass
of neutral oxygen is represented as $M_{O}$ (in $\rm M_{\odot}$), $D$ is the
SN distance (in Mpc), $F$([O I]) is the observed absolute flux of [O I] line
(in erg s-1 cm-2) and $T_{4}$ is the temperature associated with oxygen-
emitting region in units of 104 K. The [O I] 5577 / (6300, 6364) line ratio
can be used as a proxy for estimating the temperature. In our nebular phase
spectra (cf. Fig. 9), the [O I] 5577 Å line is not detected, allowing us to
set $\sim$ 0.1 as the upper limit for the line ratio. At this limit, the
emitting region should either be at a relatively low temperature ($T_{4}\leq
0.4$) for the high density limit or be at a low electron density ($n_{e}\leq$
5 $\times$ 106 cm-3) if T4 = 1 (Maeda et al., 2007a).

Previous studies have shown that during the nebular phase, the temperature of
the line emitting region lies in the range 3400 – 4200 K (Schlegel & Kirshner,
1989; Elmhamdi et al., 2004; Elmhamdi, 2011). Assuming that the high-density
regime is valid for SN 2017iro, the temperature of the line emitting region
can be taken to be $T_{4}$ = 0.4. Using the [O I] flux (in the +209 d
spectrum) of 1.12 $\times$ 10-14 erg s-1 cm-2, D = 30.8 Mpc and $T_{4}$ = 0.4,
we calculated the minimum mass of neutral oxygen to be 0.35 M⊙, which is
similar to the mass estimated for supernova iPTF13bvn (Kuncarayakti et al.,
2015b). It is to be noted that for different densities, the temperature of
line emitting regions may vary (Schlegel & Kirshner, 1989; Leibundgut et al.,
1991; Sollerman et al., 1998; Elmhamdi et al., 2004; Maeda et al., 2007a). An
increase in temperature from 4000 K to 4200 K lowers the oxygen mass by a
factor of $\sim$ 0.8. Similarly, a decrease in temperature from 4000 K to 3400
K results in higher oxygen mass by a factor of $\sim$ 3\. Hence, depending on
the assumed temperature and density, the range of oxygen mass spans 0.28 –
1.05 M⊙. The clumpy nature of the ejecta and the presence of non-optically
thin material (which is not accounted here) may raise this value (Maeda et
al., 2007b; Kuncarayakti et al., 2015b). In Table 8, the neutral oxygen mass
estimates are listed for a few well-studied objects showing a broad range of
estimated masses. We admit that the mass of the oxygen estimated here is a
lower limit of the total oxygen present in the ejecta of SN 2017iro.

Figure 15: [O I] 6300, 6364 / [Ca II] 7291, 7324 Å line ratio of CCSNe from
Kuncarayakti et al. (2015b). The [O I]/[Ca II] ratio for SN 2017iro obtained
from the nebular spectra (between +173 to +209 d) has been shown with a star
symbol. The phases (relative to the time of maximum light) are shown next to
the each data point. The left (orange) and right (grey) regions belong to Type
II and SE-SNe, respectively. The reference for spectrum are the following: SNe
1987A (Pun et al., 1995), 1992H (Clocchiatti et al., 1996), 1997D (Benetti et
al., 2001), 1998A (Pastorello et al., 2005), 1998bw (Patat et al., 2001),
1999em (Leonard et al., 2002), 2004et (Sahu et al., 2006), 2004gq (Maeda et
al., 2008b), 2005bf (Maeda et al., 2007c), 2006bp (Quimby et al., 2007),
1996cb, 2002ap, 2004fe, 2004gk, 2008ax, 2009jf (Modjaz et al., 2008b, 2014b).

###  5.4 Nature of the progenitor star

Hydrodynamic explosion models suggest that Mg and O should have similar
spatial distribution within the SN ejecta of SE-SNe (Maeda et al., 2006).
Taubenberger et al. (2009) analyzed a large number of nebular spectra of SE-
SNe and inferred that in a majority of cases, the line profiles of semi-
forbidden Mg I] 4571 Å and forbidden oxygen [O I] 6300, 6364 Å are similar,
confirming the similar distribution of Mg and O in the ejecta. The strength of
Mg I] 4571 Å line grows over time in comparison with [Ca II] 7291, 7324 Å and
[O I] because the deeper Mg-O layer of the progenitor core becomes visible
with time (Foley et al., 2003; Kuncarayakti et al., 2015b). SN 2017iro shows a
feeble presence of Mg I] 4571 Å possibly due to less stripping of its
progenitor.

The progenitor mass of a CCSN can be constrained by estimating the [O I] 6300,
6364 / [Ca II] 7291, 7324 Å emission line ratio in the nebular phase
(Kuncarayakti et al., 2015b). In contrast to Ca, the explosively synthesized O
is sensitive to the main sequence mass of the progenitor, and therefore, a
larger value implies a larger core mass (however, mixing can also play a
substantial role, see Fransson & Chevalier, 1989). Kuncarayakti et al. (2015b)
compared the [O I] / [Ca II] line ratio of SE-SNe and Type II SNe. They found
that the line ratio of Type Ib/c SNe exhibits considerable dispersion. The
observed spread in the line ratio can not be explained by the temporal
evolution of the line flux. Instead, it could arise from two different
progenitor channels for SE-SNe: a single massive star or a less massive star
in a binary system. The [O I] / [Ca II] ratio in the spectra of SN 2017iro
obtained between 173 to 209 d stays almost constant around 1. In Fig. 15, the
[O I] / [Ca II] ratio of CCSNe (from Kuncarayakti et al. (2015b)) is plotted
where SN 2017iro is shown with a star symbol. It indicates that the progenitor
of SN 2017iro is comparatively a less massive star in a binary system. A less
massive progenitor in a binary system has been proposed in several Type Ib
supernovae such as SN 2007Y (Stritzinger et al., 2009), iPTF13bvn (Bersten et
al., 2014b; Kuncarayakti et al., 2015b; Eldridge et al., 2015b). This is in
agreement with the fact that binary interaction is a favoured scenario in a
majority ($\geq$70%) of massive stars (Podsiadlowski et al., 1992; Sana et
al., 2012).

The nucleosynthesis calculations for zero-age sequence progenitor masses of
13, 15, 20, and 25 M⊙ was performed by Thielemann, Nomoto, & Hashimoto (1996),
and the corresponding oxygen mass was estimated as 0.22, 0.43, 1.48, and 3.00
M⊙, respectively. Limongi & Chieffi (2003) also computed explosive yields of
massive stars in the mass range 13 – 35 $\rm M_{\odot}$ with an initial solar
composition using the FRANEC (Chieffi et al., 1998) code. Their estimated
oxygen mass is $\sim$0.3 and $\sim$0.5 M⊙ for 13 and 15 M⊙ progenitors masses,
respectively. Both the nucleosynthesis calculations indicate that the
estimated oxygen mass of 0.35 M⊙ could be produced by a progenitor star with
zero-age main sequence mass in the range $\sim$13–15 M⊙. It is further
supported by the findings of Fang et al. (2019), where the upper limit of the
Type IIb/Ib SNe progenitors is proposed as $\sim$17 M⊙ for a progenitor in a
binary system (see, also Taddia et al., 2018; Prentice et al., 2019). It is
worth mentioning that the estimated oxygen mass for SN 2017iro (Section 5.3.2)
is only a lower limit of the total oxygen mass. With the increased uncertainty
range (0.28 – 1.05 M⊙), the range of possible progenitor masses increases
substantially, especially towards the higher progenitor masses ($\sim$20 M⊙).
However, the lower value of [O I] / [Ca II] ratio indicates towards a lower
zero-age main sequence mass progenitor of SN 2017iro.

The pre-explosion images of NGC 5480 available in the HST archive (observed on
2017 January 17, Proposal ID: 14840, PI: Andrea Bellini), were examined to
identify the possible progenitor candidate of SN 2017iro. The image was
obtained with ACS/WFC in filter F606W. A diffuse source with an extent of 5
pixels was detected in the pre-explosion images at the supernova location. At
the adopted distance of the host galaxy, the extent corresponds to $\sim$ 35
parsecs. This extended source could possibly be a star cluster hosting the
progenitor of SN 2017iro.

##  6 Summary

We have presented optical photometric (40 epochs) and spectroscopic (34
epochs) follow-up of the Type Ib SN 2017iro. The light curve evolution in
$UBVRI$ follows the usual trend of SE-SNe, i.e., blue pass-bands peak before
the red pass-bands indicating that the SN photosphere is rapidly cooling. The
$\rm\Delta m_{15}$ parameter (for $BVRI$-bands) are comparable to those
estimated for a large SE-SNe sample (Drout et al., 2011; Taddia et al., 2018)
and in particular to SN 2012au. Further, the colours of SN 2017iro evolve
similar to Type Ib events. The absolute $V$-band and bolometric
(log${}_{10}L$) luminosities are –17.76 mag and 42.39 erg s-1, respectively,
indicative of a moderately luminous SN. The explosion parameters were computed
by applying different analytical models to the quasi-bolometric light curve
which yield a 56Ni mass of $\sim$0.05 – 0.10 M⊙, $M_{\rm ej}$ of $\sim$1.4 –
4.3 M⊙ and $E_{\rm k}$ $\sim$0.8 – 1.9 $\times$ 1051 erg.

There seem to be two groups of Type Ib SNe in terms of the light curve tail
decline rates, suggesting a heterogeneity (Wheeler et al., 2015) in the late
time light curve evolution. However, lack of an adequate amount of data during
the late phase of these events could be a possible reason for such inference.
Hence, late phase monitoring of a larger sample of such events may provide a
clearer insight to this finding. The late phase light curves of SN 2017iro
decayed with moderate steepness, indicating incomplete $\gamma$-ray trapping
by the SN ejecta. A low to moderate level of 56Ni mixing is supported by the
non-monotonic evolution of colour curves (Dessart et al., 2012; Yoon et al.,
2019) in the case of SN 2017iro. However, we can not draw a firm conclusion
due to unavailability of sufficient early phase data of this event.

In the first optical spectrum ($-7$ d) of SN 2017iro, the He i 5876 Å feature
is clearly identified. Other He i lines such as 4471, 6678, and 7065 Å are
comparatively weak in the beginning, but they never became as strong as He i
5876 Å. The ejecta velocity estimated using the Fe ii 5169 Å line is $\sim$
9000 km s-1 near $B$-band maximum which became almost constant to $\sim$ 5500
km s-1 beyond +25 d. The Mg i] line is very weak during the nebular phase in
SN 2017iro. The general spectroscopic behaviour of SN 2017iro is similar to SN
2009jf and iPTF13bvn. The metallicity near SN 2017iro location (12 + log (O/H)
= 8.64) is similar to the solar value. A neutral oxygen mass of $\sim$ 0.35 M⊙
was estimated in the ejecta. The [O i] / [Ca ii] line ratio observed in SN
2017iro indicates the progenitor had a main sequence mass of $\sim$13 – 15 M⊙.
The inferred neutral oxygen mass and the progenitor main sequence mass for SN
2017iro are very similar to the type Ib supernova iPTF13bvn that occurred in a
binary system.

## Acknowledgments

We thank the referee for critical review and constructive suggestions that
helped to improve the paper. We also thank Sung-Chul Yoon and Hanindyo
Kuncarayakti for sharing the data. BK, DKS, and GCA acknowledge BRICS grant
DST/IMRCD/BRICS/PilotCall1/MuMeSTU/2017(G) for the present work. DKS and GCA
also acknowledge DST/JSPS grant, DST/INT/JSPS/P/281/2018. We are grateful to
the observers at HCT who provided their valuable time to monitor this event
and also thank the staff of IAO, Hanle, and CREST, Hosakote, that made these
observations possible. The Weizmann interactive supernova data repository
(WISeREP) - http://wiserep.weizmann.ac.il is also acknowledged. This research
has made use of the NASA/IPAC Extragalactic Database (NED), which is operated
by the Jet Propulsion Laboratory, California Institute of Technology, under
contract with the National Aeronautics and Space Administration. This research
made use of RedPipe222https://github.com/sPaMFouR/RedPipe (Singh, 2021), an
assemblage of data reduction and analysis scripts written by AS.

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