MNRAS 000, 1–24 (2022) Preprint 27 October 2022 Compiled using MNRAS LATEX style file v3.0
SN 2020wnt: a slow-evolving carbon-rich superluminous supernova
with no O ii lines and a bumpy light curve
C. P. Gutiérrez ,1,2,★, A. Pastorello, ,3 M. Bersten,4,5,6 S. Benetti,3 M. Orellana,7,8 A. Fiore,9,10,3
E. Karamehmetoglu, ,11 T. Kravtsov,2 A. Reguitti, ,12,13,3 T. M. Reynolds,14,15 G. Valerin,16,3
P. Mazzali,17,18 M. Sullivan, ,19 Y.-Z. Cai,20 N. Elias-Rosa, ,3,21 M. Fraser, ,22, E. Y. Hsiao, ,23
E. Kankare, ,2,24, R. Kotak,2 H. Kuncarayakti,2,1 Z. Li,25,26 S. Mattila,2 J. Mo,27 S. Moran,2
P. Ochner,3,16 M. Shahbandeh, ,23 L. Tomasella, ,3 X. Wang,20,27 S. Yan,20 J. Zhang,28,29 T. Zhang,25,26
M. D. Stritzinger ,11
Affiliations are listed at end of paper
Accepted XXX. Received YYY; in original form ZZZ
ABSTRACT
We present the analysis of SN 2020wnt, an unusual hydrogen-poor super-luminous supernova
(SLSN-I), at a redshift of 0.032. The light curves of SN 2020wnt are characterised by an
early bump lasting ∼ 5 days, followed by a bright main peak. The SN reaches a peak absolute
magnitude of M𝑚𝑎𝑥𝑟 = −20.52 ± 0.03 mag at ∼ 77.5 days from explosion. This magnitude
is at the lower end of the luminosity distribution of SLSNe-I, but the rise-time is one of the
longest reported to date. Unlike other SLSNe-I, the spectra of SN 2020wnt do not show O ii,
but strong lines of C ii and Si ii are detected. Spectroscopically, SN 2020wnt resembles the
Type Ic SN 2007gr, but its evolution is significantly slower. Comparing the bolometric light
curve to hydrodynamical models, we find that SN 2020wnt luminosity can be explained by
radioactive powering. The progenitor of SN 2020wnt is likely a massive and extended star
with a pre-SN mass of 80 M and a pre-SN radius of 15 R that experiences a very energetic
explosion of 45 × 1051 erg, producing 4 M of 56Ni. In this framework, the first peak results
from a post-shock cooling phase for an extended progenitor, and the luminous main peak is
due to a large nickel production. These characteristics are compatible with the pair-instability
SN scenario. We note, however, that a significant contribution of interaction with circumstellar
material cannot be ruled out.
Key words: supernovae: general – supernovae: individual: SN 2020wnt
1 INTRODUCTION
The rise of wide-field sky surveys in the last decade revealed the
existence of very bright supernovae (SNe), now known as super-
luminous supernovae (SLSNe; Gal-Yam 2012). SLSNe are around
two orders of magnitude brighter than classical SNe (𝑀𝑉 < −19.5
mag; Quimby et al. 2011; Gal-Yam 2012; Angus et al. 2019; Gal-
Yam 2019; Inserra 2019; Nicholl 2021), and show a large diversity
in both their light curves and spectra. Their host galaxies are gen-
erally found to be faint dwarf galaxies (Neill et al. 2011) with low
metallicity (< 0.5 Z) and low stellar masses (e.g. Stoll et al. 2011;
Lunnan et al. 2014; Perley et al. 2016; Chen et al. 2017; Schulze
et al. 2018).
Initially, SLSNe were sub-classified into two classes:
hydrogen-poor (SLSN-I) and hydrogen-rich events (SLSNe-II; Gal-
★ E-mail: claudia.gutierrez@utu.fi
Yam 2012). However, with the increase in the number of objects, es-
pecially thosewith better data sets, more detailed sub-classifications
have been necessary.Within the SLSNe-I class, it is possible to iden-
tify two sub-groups based on their distinctive photometric proper-
ties: the slow-evolving SLSNe-I that show long rise times (> 50
days) to the main peak, and the fast-evolving SLSNe-I that have
rise times shorter than ∼ 30 days (Inserra et al. 2017; Quimby et al.
2018; Inserra 2019). Furthermore, Inserra et al. (2018) found that
slow-evolving SLSNe-I have small expansion velocities (𝑣 . 10000
km s−1) and almost non-existent velocity gradients (Δ𝑣/Δ𝑡 in units
of km s−1day−1, over the time interval [+10, +30]), while the fast-
evolving subgroup members have large velocities and large velocity
gradients. However, the identification of SLSNe-I with intermedi-
ate (or transitional) properties (e.g. Gaia16apd; Kangas et al. 2017;
SN 2017gci; Fiore et al. 2021), suggests a continuum distribution
(Nicholl et al. 2015a; De Cia et al. 2018). Therefore, the bimodality
© 2022 The Authors
ar
X
iv
:2
20
6.
01
66
2v
2 
 [a
str
o-
ph
.H
E]
  2
6 O
ct 
20
22
2 Gutiérrez et al.
or separation found by Inserra et al. (2018) could be the consequence
of the small sample considered (only 18 SNe).
Studies of single objects with good photometric and spectro-
scopic coverage have revealed a large number of unusual properties
that can give insights on the progenitor and explosion mechanisms.
For instance, observations revealed pre-peak bump light curve mor-
phologies (Leloudas et al. 2012; Nicholl et al. 2015b; Smith et al.
2016; Anderson et al. 2018; Angus et al. 2019), and light curve un-
dulations (Nicholl et al. 2015b; Inserra et al. 2017; Yan et al. 2017a;
Fiore et al. 2021). To explain the pre-peak bump, at least threemech-
anisms have been proposed. These include shock breakout within a
dense circumstellar material (CSM; Moriya & Maeda 2012), shock
cooling of extended material (e.g. Nicholl et al. 2015b; Piro 2015;
Smith et al. 2016;Vreeswijk et al. 2017), and an enhancedmagnetar-
driven shock breakout (Kasen et al. 2016). On the other hand, the
light curve undulations have been interpreted as a signature of the
interaction of the ejecta with CSM (Gal-Yam et al. 2009; Inserra
et al. 2017; Inserra 2019; but see, Moriya et al. 2022; Kaplan &
Soker 2020).
Spectroscopically, the W-shaped O ii features are a key char-
acteristic of SLSNe-I and have been recognised as such since their
discovery (Quimby et al. 2011; Mazzali et al. 2016), although re-
cently, it has been found that SLSN-I can be separated into two
different sub-classes based on their pre-maximum spectra: events
showing the W-shaped O ii features, and events which do not show
such W-shaped absorption (e.g. Könyves-Tóth & Vinkó 2021). At
early times, SLSN-I spectra are also characterised by the presence
of C ii (e.g. Dessart et al. 2012; Mazzali et al. 2016; Dessart 2019;
Gal-Yam 2019) and Si ii (Inserra et al. 2013). Despite the limited
number of late-time spectra available for SLSNe-I, they appear to re-
semble SNe Ic associated with gamma-ray bursts (e.g. SN 1998bw;
Nicholl et al. 2016b; Jerkstrand et al. 2017). The similarities between
SLSNe-I and SNe Ic suggest they are somewhat related (Pastorello
et al. 2010).
Diverse explosion scenarios have been proposed to explain
SLSNe (see review of Moriya et al. 2018, and references therein).
These include pair-instability mechanism (Heger & Woosley 2002;
Gal-Yam et al. 2009), the interaction of the SN ejecta with CSM
(e.g. Chevalier & Irwin 2011; Chatzopoulos et al. 2011; Ginzburg
& Balberg 2012; Dessart et al. 2015; Sorokina et al. 2016), and the
spin-down of a rapidly rotating, highlymagnetic neutron star (Kasen
2010; Woosley 2010; Bersten et al. 2016). Stars with initial masses
larger than 140 M are predicted to undergo pair instability and
explode completely (Barkat et al. 1967; Rakavy & Shaviv 1967).
The light curves of these pair-instability supernovae (PISNe) are
expected to be very luminous, therefore, they have been proposed
as a good alternative to explain the high luminosity, and in turn,
the large amounts of synthesised 56Ni in SLSNe. However, the light
curves and spectra of some observed objects are not compatible
with this scenario (e.g. Dessart et al. 2013; Jerkstrand et al. 2016;
Mazzali et al. 2019). Another scenario is the interaction of the SN
ejecta with CSM produced by mass-loss of the progenitor star prior
the explosion. This mechanism offers a proper explanation for a
luminous and bumpy (fluctuations in brightness) light curves, but
the absence of narrow emission lines in the SN spectra is currently
a major issue (but see, Chevalier & Irwin 2011). The most accepted
alternative of powering source of many SLSNe I has been found
in the magnetar scenario, which can explain most of the observed
properties in SLSNe. However, Soker &Gilkis (2017) found that the
energy of the explosion in the magnetar model are more than what
the neutrino-driven explosion mechanism can supply, therefore, a
jet feedback mechanism from jets launched at magnetar birth may
be involved.
Given that several open questions remain regarding both the ex-
plosion mechanism and the progenitors of SLSNe, studying nearby
SLSNe in detail allows us to discriminate among various scenarios.
In this paper, we present SN 2020wnt, one of the closest (z=0.032)
SLSNe-I discovered to date. The excellent coverage from explo-
sion to ∼ 500 days allows us to characterise its properties. Its light
curves show an early bump, followed by a slow rise to the main
peak. Some fluctuations in brightness are also observed at late time.
On the other hand, unlike other SLSN objects, SN 2020wnt spectra
do not show signs of O ii, but its evolution resembles that of the
type Ic carbon-rich SN 2007gr (Valenti et al. 2008a; Hunter et al.
2009; Chen et al. 2014). This unusual similarity provides an ex-
cellent opportunity to understand the possible connection between
H-poor SNe and SLSNe.
The paper is organised as follows. A description of the obser-
vations and data reduction is presented in Section 2. The character-
isation of SN 2020wnt (host galaxy, photometric and spectral prop-
erties) is given in Section 3. In Section 4, comparisons with similar
objects are presented, while in Section 5 the explosion progenitor
properties are analysed through hydrodynamical modelling. Finally,
in Section 6 and Section 7, we present the discussion and conclu-
sions, respectively. Throughout this work, we will assume a flat
ΛCDMuniverse, with aHubble constant of𝐻0 = 70 km s−1Mpc−1,
and Ωm = 0.3.
2 OBSERVATIONS OF SN 2020WNT
2.1 Detection and classification
SN 2020wnt (a.k.a. ZTF20acjeflr and ATLAS20beko) was detected
by the Zwicky Transient Facility (ZTF; Bellm et al. 2019; Graham
et al. 2019) on 2020 October 14 (MJD= 59136.40), at a magnitude
m𝑔 = 19.70±0.11mag. A couple of hours later (MJD= 59136.47),
a detection in the 𝑟−band confirmed the new object (m𝑟 = 19.57 ±
0.05mag). The discoverywas reported to the TransientNameServer
(TNS1) by the Automatic Learning for the Rapid Classification of
Events (ALeRCE) broker (Förster et al. 2021) on MJD= 59136.79.
SN 2020wnt was also detected by the Asteroid Terrestrial-impact
Last Alert System (ATLAS; Tonry et al. 2018; Smith et al. 2020)
on 2020 October 14 (MJD= 59136.50; m𝑐 = 19.70 ± 0.11 mag).
The last non-detection obtained by ZTF was on 2020 October 12
(MJD= 59134.45) with a detection limit of m𝑟 ∼ 20.70 mag. A
deeper non-detection in the 𝑔 band (∼ 21.00 mag) occurred earlier
the same night (MJD= 59134.39). Using these constraints, we adopt
the mid-point between the last non-detection and first detection as
the explosion epoch (MJD= 59135.42 ± 0.98; 2020 October 13).
SN 2020wnt was spectroscopically observed on 2020 November 15
(MJD= 59168.0) by the UC Santa Cruz group and classified as a
SN I at a redshift of 0.032 (Tinyanont et al. 2020).
2.2 Photometry
SN 2020wnt was observed photometrically for 72 weeks, from 2020
October 14 to 2022 February 27, using various facilities.Most of the
observations were carried out by two wide-field imaging surveys,
namely ATLAS and ZTF. From 2020 October 14 to 2021 November
4, photometry in the orange (𝑜) filter (a red filter that covers a
1 https://wis-tns.weizmann.ac.i
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 3
wavelength range of 5600 to 8200 Å) and cyan (𝑐) filter (wavelength
range 4200 to 6500 Å) was obtained by the twin 0.5 m ATLAS
telescope system (Tonry et al. 2018). ATLAS photometry (Tonry
et al. 2018 and Smith et al. 2020) was obtained through the ATLAS
forced photometry server2. ZTF obtained 𝑔− and 𝑟−band images
from 2020 October 14 to 2021 November 9. The ZTF photometry
was obtained through the ZTF forced-photometry service (Masci
et al. 2019). The light curves were generated following the steps
presented in the ZTF documentation3.
Optical imaging was obtained with the Copernico 1.82 m tele-
scope equipped with Asiago Faint Object Spectrograph and Camera
(AFOSC) and the 67/91 Schmidt Telescope equippedwithMoravian
G4-16000LC at the Asiago Observatory (Italy); the 0.8-m Tsinghua
University–NAOC (National Astronomical Observatories of China)
Telescope (TNT) at Xinglong Observatory of NAOC (Huang et al.
2012); the 2-m Liverpool Telescope (LT) using the IO:O imager,
the Low-Resolution Spectrograph (LRS) at the 3.6-m Telescopio
Nazionale Galileo (TNG), and the 2.56-mNordic Optical Telescope
(NOT) using the Alhambra Faint Object Spectrograph and Camera
(ALFOSC) at the Roque de LosMuchachosObservatory (Spain). 10
epochs of near-infrared (NIR; 𝐽𝐻𝐾) photometrywere obtainedwith
NOTCam at NOT, while six epochs of UltraViolet (UV) andOptical
observations were obtained with the UltraViolet/Optical Telescope
(UVOT) onboard the Neil Gehrels Swift Observatory spacecraft.
All NOT observations were obtained through the NOT Unbiased
Transient Survey 2 (NUTS24) allocated time.
Data reduction and SN photometry measurements for Asiago,
LT and NOT were performed using the python/pyraf SNOoPY
pipeline (Cappellaro 2014), whereas the TNT andTNG imageswere
reduced with iraf following standard procedures. The photometry
for TNT and TNG was performed using the python package pho-
tutils (Bradley et al. 2019) of astropy (Astropy Collaboration
et al. 2018). All 𝑢𝑔𝑟𝑖𝑧 magnitudes were calibrated using observa-
tions of local Sloan and Pan-STARRS sequences (Chambers et al.
2016; Magnier et al. 2020). The 𝐵𝑉 magnitudes were derived using
Pan-STARRS and the transformations in Chonis & Gaskell (2008),
while the 𝐽𝐻𝐾 magnitudes were calibrated using 2MASS (Skrut-
skie et al. 2006). UVOT reductions and the resulting photometry
were performed by using the heasoft Software5 (NasaHigh Energy
Astrophysics Science Archive Research Center (Heasarc) 2014) and
aperture photometry.
Optical (𝑢𝐵𝑔𝑉𝑟𝑖𝑧), NIR and UVOT photometry are presented
in Tables A1, A2 and A3, respectively. The mean magnitudes from
ATLAS are listed in Table A4, while the 𝑔𝑟 photometry from ZTF
is in Table A5.
2.3 Spectroscopy
26 optical spectra of SN 2020wnt were obtained spanning phases
between 32 and 293 days from explosion. These observations were
acquired with seven different instruments: ALFOSC at the Nordic
Optical Telescope (NOT), AFOSC at the Copernico 1.82-m Tele-
scope (Mount Ekar); LRS at the Nazionale Galileo (TNG), Yunnan
Faint Object Spectrograph and Camera (YFOSC) at the Lijiang 2.4-
m Telescope (LJT), Beijing Faint Object Spectrograph and Camera
2 https://fallingstar-data.com/forcedphot/
3 https://irsa.ipac.caltech.edu/data/ZTF/docs/forcedphot.
pdf
4 https://nuts.sn.ie
5 http://heasarc.gsfc.nasa.gov/ftools
3h46m44.0s 40.0s 36.0s 32.0s
43◦15’00”
14’00”
13’00”
RA
D
E
C
SN 2020wnt
E
N
r− band
WISEA J034638.04+431348.3
Figure 1. NOT 𝑟−band image of SN 2020wnt and its host galaxy, WISEA
J034638.04+431348.3 The red circle marks the SN (RA = 03h46m37.s95
Dec = +43◦13′45.′′30 (J2000)), while the blue ellipse marks its host galaxy.
The orientation of the image is indicated in the bottom-right corner.
(BFOSC) at the Xinglong 2.16-m Telescope (XLT), Optical Sys-
tem for Imaging and low-Intermediate-Resolution Integrated Spec-
troscopy (OSIRIS) at the 10.4-m Gran Telescopio Canarias (GTC),
and Kast double spectrograph on the 3.0-m Shane telescope at the
Lick Observatory6. All spectra were reduced using standard iraf
routines (bias subtraction, flat-field correction, 1D extraction, and
wavelength calibration). The flux calibration was performed using
spectra of standard stars obtained during the same night. For theAL-
FOSC, AFOSC and OSIRIS spectra, the data were reduced using
the foscgui7 pipeline.
Additionally, a NIR spectrum was obtained at ∼ 49 days from
explosion with the 0.7-5.3 Micron Medium-Resolution Spectro-
graph and Imager (SpeX instrument) on the 3.2-m NASA Infrared
Telescope Facility. The spectrum was taken in cross-dispersed SXD
mode with the 0.5 arcsec slit, and reduced using the spextool
software package (Cushing et al. 2004) following the prescriptions
described by Hsiao et al. (2019). Details of the instruments used for
the spectroscopic observations are reported in Table A6. All spectra
will be available through theWISeREP8 archive (Yaron &Gal-Yam
2012).
3 CHARACTERISING SN 2020WNT
3.1 Host galaxy
The host galaxy of SN 2020wnt isWISEA J034638.04+431348.3, a
faint galaxy with no published redshift or distance information. The
redshift adopted in our analysis is derived from the narrow emission
lines (H𝛼, [O iii] 𝜆5007) visible in the SN spectrum. These lines
give us a mean redshift of 0.032. Given the lack of independent
6 Public spectrum obtained from the TNS webpage: https://www.
wis-tns.org/object/2020wnt
7 FOSCGUI is a graphical user interface aimed at extracting SN spec-
troscopy and photometry obtained with FOSC-like instruments. It was
developed by E. Cappellaro. A package description can be found at
sngroup.oapd.inaf.it/foscgui.html.
8 http://wiserep.weizmann.ac.il/home
MNRAS 000, 1–24 (2022)
4 Gutiérrez et al.
measurements of distance to this galaxy, we estimate the uncertainty
in our measurements assuming a peculiar velocity of ±200 km s−1
(Tully et al. 2013). With these values, we compute a distance of
𝑑 = 140.4 ± 3.0Mpc, which corresponds to a distance modulus of
𝜇 = 35.74 ± 0.05 mag. Figure 1 shows the NOT 𝑟−band image of
SN 2020wnt and its host galaxy.
TheGalactic reddening is quite highwith 𝐸 (𝐵−𝑉) = 0.42mag
(Schlafly & Finkbeiner 2011), while the host galaxy component is
negligible. This is determined by the absence of narrow interstellar
Na iD lines (𝜆𝜆5889, 5895) at the rest wavelength of the host. As
shown in section 3.4, the spectra of SN 2020wnt display a strong
Na iD line, but it corresponds to the Milky Way component. There-
fore, we assume that the reddening in the direction of SN 2020wnt
is totally dominated by the Milky Way.
To characterise the global properties of the SN 2020wnt host
galaxy, we got 𝑔𝑟𝑖𝑧 photometry from the Pan-STARRS1 (PS1) pub-
lic science archive9. To obtain estimates of the stellar mass (𝑀∗)
and star formation rate (SFR), we use a custom galaxy spectral
energy distribution (SED) fitting code, following the procedure de-
tailed in Sullivan et al. (2010). The code is similar to z-peg (Le
Borgne & Rocca-Volmerange 2002), but uses the stellar population
templates of pégase.2 (Fioc & Rocca-Volmerange 1997). The best-
fitting templates correspond to 𝑀∗ = log (𝑀/𝑀) = 8.22+0.20−0.12 and
log (SFR) = −4.89+3.09−5.11 Myr−1. The large uncertainties obtained
for the SFR are mainly associated to the lack of photometry in bluer
bands (Childress et al. 2013). Such low stellar mass and star for-
mation obtained for the host galaxy of SN 2020wnt are consistent
with the expected range measured for other SLSN hosts (Neill et al.
2011; Chen et al. 2013; Lunnan et al. 2014; Leloudas et al. 2015;
Perley et al. 2016; Schulze et al. 2018). In absence of a host spec-
trum, we could infer the metallicity through the mass – metallicity
relation (e.g Tremonti et al. 2004). Following the prescriptions of
Kewley & Ellison (2008), we derive a value that corresponds to 12
+ log(O/H)= 8.17 ± 0.11 dex in the O3N2 calibration and and 12
+ log(O/H)= 8.18 ± 0.09 dex in N2 calibration, respectively. These
values suggest a low metallicity and are consistent with previous
findings for SLSNe (e.g., Chen et al. 2017).
3.2 Light curves
Figure 2 (top panel) shows the rest-frame multiband light curves of
SN 2020wnt. The excellent photometric coverage during the first ∼
320 days fromexplosion allows us to constrain exceptionallywell the
light curve shape, magnitudes at maximum, rise times, and decline
rates in the different bands. To estimate the main parameters of the
light curves, we use Gaussian processes (GPs). For this procedure,
we use the python package george (Ambikasaran et al. 2016),
following the prescriptions of Gutiérrez et al. (2020a).
As seen in the top panel of Figure 2, SN 2020wnt evolves quite
slowly and shows several distinctive properties during its evolution.
The dense sampling in three of the four filters with very early
observations (𝑔𝑐𝑟) allows us to detect an initial peak, which is
brighter in the bluer filters with an absolute magnitude of M𝑔 =
−17.71 mag, M𝑐 = −17.38 mag, M𝑟 = −17.30 mag.
Following this initial peak, the 𝑔𝑐𝑟 light curves show a decrease
in brightness (between 0.5 in 𝑟 and 0.7 mag in 𝑔). At this point,
the SN reaches a minimum value, with absolute magnitudes of
M𝑔 = −17.00 mag, M𝑐 = −17.02 mag, and M𝑟 = −16.8 mag at
∼ 4.9 days. After this phase, a rise of ∼ 3.3 − 3.7 mag is observed
9 https://catalogs.mast.stsci.edu/panstarrs/
Table 1. Light curve parameters of SN 2020wnt.
Peak Abs. Mag Rise time ΔM(Peak−130) ΔM(77.5−130)
Band (mag)★ (days)★ (mag)† (mag)
(1) (2) (3) (4) (5)
𝑢 −20.00 ± 0.02 61.3 ± 1.4 3.58 3.25
𝐵 −20.27 ± 0.01 65.0 ± 2.2 2.49 2.45
𝑔 −20.52 ± 0.03 72.0 ± 2.5 2.27 2.25
𝑐 −20.50 ± 0.03 70.2 ± 1.4 1.67 1.62
𝑉 −20.50 ± 0.02 72.4 ± 3.1 1.69 1.65
𝑟 −20.52 ± 0.03 77.5 ± 3.1 1.32 1.32
𝑜 −20.49 ± 0.01 75.0 ± 1.8 1.29 1.29
𝑖 −20.40 ± 0.02 76.8 ± 5.1 1.09 1.08
𝑧 −20.16 ± 0.01 79.0 ± 4.1 0.81 0.81
𝐽 −20.84 ± 0.10 89.0 ± 5.2 1.01 0.94
𝐻 −20.93 ± 0.08 93.8 ± 5.1 0.72 0.55
𝐾 −20.92 ± 0.03 101.7 ± 1.0‡ 0.71 –
Columns: (1) Band; (2) Peak absolutemagnitudes; (3) Rise time; (4) Change
in magnitude from the peak to 130 d from explosion; (5) Change in magni-
tude from 77.5 to 130 d from explosion.
★ Peak absolute magnitudes and rise times were obtained from Gaussian
Process (GP) fits. Magnitudes are corrected by the Milky Way extinction.
† Change in magnitude from the peak. Note that each band reaches the peak
at different times.
‡ Due to the lack of data in 𝐾 , the first point is assumed as the maximum
and the uncertainty of the maximum time is the error from explosion epoch.
in all bands. A high cadence follow-up in the 𝑢𝐵𝑔𝑉𝑟𝑖𝑧 bands starts
after about 38 days from the explosion. This permits to cover the
SN maximum in all optical bands.
From the GP fits, we find that SN 2020wnt reaches the main
peak in the optical bands between ∼ 61 and 79 days from explosion
(all phases stated in this paper are in the rest frame). The rise
times are different in each band, with a faster rise in 𝑢 and a more
extended rise in the NIR bands. The absolute magnitude peaks
are around −20.00 mag in 𝑢 and −20.93 mag in 𝐻. In 𝑔𝑐𝑉𝑟𝑜 the
absolute peak magnitudes are around −20.50 mag. These values
place SN 2020wnt at the bottom of the luminosity distribution of
SLSNe-I (e.g. Angus et al. 2019). Table 1 shows the absolute peak
magnitudes and the rest-frame rise times obtained in all optical and
NIR bands.
One interesting characteristic observed in the light curves of
SN 2020wnt is the behaviour around peak. In the redder bands, the
SN reaches the maximum a bit later than in the blue bands, and
the luminosity stays quasi-constant for a longer time, displaying a
kind of "plateau". After maximum, the decline in the blue bands
is much faster than in the redder ones. From the main peak to 130
days, the SN dims by ∼ 3.58 mag in the 𝑢 band, 2.27 mag in 𝑔 and
1.32 mag in 𝑟 (almost three times slower than 𝑢). As the 𝑢 peak
occurs before than the peak in 𝑔 and 𝑟 (61.3, 72.0, and 77.5 days
in 𝑢, 𝑔 and 𝑟 , respectively), we can better compare these decreases
by fixing a range of time. Measuring the change in magnitude from
77.5 days (the epoch of 𝑟−band peak) to 130 days, we see that the
SN dims by ∼ 3.26 mag in 𝑢 and 2.25 mag in 𝑔. The decline rates
in these two ranges are also presented in Table 1. Moreover, in the
𝑈𝑉𝑊1 and𝑈𝑉𝑊2 bands, we see a flattening or upturn, which may
be attributable to the known red leak10. Meanwhile, for the𝑈𝑉𝑀2
10 SWIFT 𝑈𝑉𝑊 1 and 𝑈𝑉𝑊 2 filters have extended red tails that reach
into the optical. When the flux is optically dominated, the UV contri-
bution can be minimal (https://swift.gsfc.nasa.gov/analysis/
uvot_digest/redleak.html)
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 5
10
15
20
25
30
A
p
p
ar
en
t
M
ag
n
it
u
d
e
K − 5.5
H − 4.5
J − 3.5
z − 3
i− 2
o− 1
r
V + 1
c+ 2
g + 3
B + 4
u+ 5
UVW1 + 6
UVM2 + 7
UVW2 + 8
−25
−20
−15
−10
−5
A
b
so
lu
te
M
ag
n
it
u
d
e
−100 0 100 200 300 400
Rest-frame days from r−band maximum light
AFOSC
ALFOSC
ATLAS
IO:O
LRS
Moravian CCD
NOTCAM
Swift
TNT
ZTF 0 20
Rest-frame days from expl.
18
20
22A
p
a
re
n
t
m
a
g
−100 0 100 200 300 400 500
Rest-frame days from explosion
0
1
In
st
ri
nt
ic
C
ol
ou
r B − V
g − r
r − i
i− z
Figure 2. Top: UV, optical and NIR light curves of SN 2020wnt. Upper limits are presented as open symbols. The explosion time is indicated as a vertical
dashed line. The photometry is corrected for Milky Way extinction using the Cardelli et al. (1989) extinction law. The solid lines show the Gaussian process
(GP) interpolation before 180 days, and the shaded regions represent the errors from the GP. UV and 𝐵𝑉 𝐽𝐻𝐾 photometry are in the Vega system, while
𝑢𝑔𝑐𝑟𝑜𝑖𝑧 photometry is in the AB system. The inset plot shows the light curves in 𝑔𝑐𝑟 at very early phases. Bottom: Intrinsic colour curves of SN 2020wnt.
filter, which is not affected by the red leak, we only measure upper
limits. Therefore, the evolution of the Swift UV bands indicates a
decline in the UV flux and the emergence of an optically dominated
spectral energy distribution (SED) ∼100 days post explosion.
After ∼ 130 days, the drop in brightness slows down. Fitting
a line to the observations obtained between 130 and 180 days (the
last observation before the SN went behind the sun), we measure
a slope of 1.01 ± 0.03 mag per 100 days in 𝑟, and 1.18 ± 0.01
per 100 days in 𝑖. The monitoring of SN 2020wnt restarted at 247
days in 𝑟, and a couple of days later in 𝐵𝑔𝑉𝑐𝑖𝑧𝐽𝐻𝐾 . With the new
data, we again fit a line to the data between 130 and 275 days and
we found slower declines, with slopes of 0.90 ± 0.01 mag per 100
days in 𝑟 and 0.99 ± 0.06 mag per 100 days in 𝑖. These values are
very close to those expected from the 56Co decay (0.98 mag per
100 days; Woosley et al. 1989). Starting from 273 days, the SN is
found to experience a drastic and sudden drop in brightness in all
bands. Within ∼ 35 days, the magnitudes decrease by about 1.5
mag. Fitting a line after 273 days in the 𝑟−band, we find a slope
of 4.5 ± 0.3 mag per 100 days. From ∼ 320 days, fluctuations in
brightness are observed both in the optical and the NIR bands.
The intrinsic colour curves of SN 2020wnt are presented in
the bottom panel of Figure 2. During the first ∼ 35 - 40 days, we
can only infer 𝑔 − 𝑟 colour information. In ∼ 20 days, SN 2020wnt
becomes redder, going from a 𝑔 − 𝑟 = −0.41 mag to 𝑔 − 𝑟 = 0.16
MNRAS 000, 1–24 (2022)
6 Gutiérrez et al.
41.5
42.0
42.5
43.0
43.5
lo
g 1
0
L
bo
l
(e
rg
s−
1
)
56Co decay
SN 2020wnt
0 100 200
Rest-frame days from Lbol maximum
5
10
T
B
B
[1
03
K
]
0 100 200 300
Rest-frame days from explosion
0.0
2.5
5.0
7.5
R
ad
iu
s
[1
01
5
cm
]
Figure 3. Top: Bolometric light curve of SN 2020wnt. The dashed red
line shows the luminosity expected from 56Co decay (assuming full trap-
ping). Middle: Temperature evolution of SN 2020wnt. Bottom: Evolution
of blackbody radius of SN 2020wnt. Error bars are comparable to the size
of the symbols.
mag. The 𝑔 − 𝑟 colour shows an initial peak at 21.3 days. After this,
the SN goes back to bluer colours, reaching a value of 𝑔−𝑟 = −0.09
mag at 42.5 days. From this epoch, the 𝑔 − 𝑟 colour shows a quasi
flat evolution up to ∼ 70 days. Later than 70 days, the SN becomes
redder again, reaching its main peak at 135.4 days with a colour
of 𝑔 − 𝑟 = 1.10 mag. A gap in the observations prevented us from
monitoring the SN evolution between 151 and 251 days, however,
when the SN is recovered, we measure a colour of 𝑔− 𝑟 = 0.10mag
showing that SN 2020wnt became bluer again, but after this, it gets
redder one more time.
Starting from ∼ 38 days, we also have the 𝐵 − 𝑉 , 𝑟 − 𝑖 and
𝑖 − 𝑧 colours. 𝐵−𝑉 shows a similar behaviour than that observed in
𝑔 − 𝑟 , but with redder colours at all phases. Unlike 𝑔 − 𝑟 and 𝐵 −𝑉 ,
the evolution of 𝑟 − 𝑖 and 𝑖 − 𝑧 shows little variation. Overall, they
tend to be bluer up to ∼ 150 days. Of these two, the colour change
is more significant in 𝑖 − 𝑧, going from −0.22 to 0.12mag (at ∼ 150
days) in comparison to the evolution from −0.13 to 0.16 mag in
𝑟 − 𝑖. After the gap, the SN gets bluer in 𝐵 −𝑉 , 𝑟 − 𝑖 and 𝑖 − 𝑧. This
tendency is clear until 320 days. From there, the temporal coverage
does not allow us to estimate a trend.
3.3 Bolometric light curve
To build the bolometric light curve of SN 2020wnt using our red-
dening corrected photometry from UV to NIR bands, we employ
the superbol code11 (Nicholl 2018). In order to have a similar cov-
erage in the different bands at the same epochs, we interpolate and
extrapolate the light curves assuming constant colours and using
the 𝑟−band as a reference filter from explosion to 320 days after
the explosion. We converted all magnitudes to flux and construct
the SED at all epochs. We computed multiple pseudo-bolometric
light curves by performing trapezoidal integration just in the optical
and NIR bands, and UV + optical + NIR. We also calculated a full
bolometric light curve by fitting a blackbody to the SED. When
comparing the bolometric light curve from UV to NIR with that
obtained by extrapolating the SED constructed from the optical and
NIR bands, we find that they are consistent. The bolometric light
curve is presented in the top panel of Figure 3.
Following the process described in Section 3.2, we use a GP to
estimate the main parameters in the bolometric light curve. We find
a peak luminosity of 𝐿𝑏𝑜𝑙 = 4.25(±0.30) ×1043 erg s−1 at 65 days.
This maximum occurs earlier than most of the peaks obtained from
the optical and NIR bands, except for the 𝑢𝐵 filters. Between 140
and 270 days, the light curve declines at a rate of 0.77 ± 0.02 mag
per 100 days. After that, we estimate a decline rate of 4.07 ± 0.01
mag per 100 days.
From superbol, the blackbody temperature (T𝐵𝐵) and radius
(R𝐵𝐵) are also obtained by fitting the SED of each epoch with a
blackbody function. Figure 3 (middle and bottom panels) shows the
temperature and radius evolution from explosion to ∼ 150 days.
At early times the temperature is relatively low, with a value of
T𝐵𝐵 ≈ 8000 K. The temperature increases and reaches a value of
T𝐵𝐵 ≈ 10000K at ∼ 51 days from explosion, and then it decreases.
On the other hand, the radius shows a continuous increase up to 105
days, where it reaches its maximum value (R𝐵𝐵 = 6.2 × 1015 cm).
Fitting a line between explosion and the maximum value, we find a
slope of ∼ 6000 km s−1. After this peak, the radius shows a slow
and steady decline to 5.5 × 1015 cm.
3.4 Spectral evolution
Figure 4 shows the spectral evolution of SN 2020wnt covering the
phases from 32 to 293 days after explosion. The slow evolution
detected in the light curves is also visible in the spectra, where a
blue continuum is observed for around 100 days. The first spectrum
taken at 32 days (−46 d from the maximum light in the 𝑟 band) and
used for the classification, is dominated by strong lines of O i 𝜆7774,
Ca ii (H&K and NIR triplet), Si ii 𝜆6355, and C ii 𝜆6580, 𝜆7235.
Na iD /He i and the Fe ii 𝜆4924, 5018, 5169 lines are also clearly
detected. From 32 to 78 days after explosion, there are limited
changes in the spectra, with small variations in the relative line
intensities. More precisely, C i 𝜆6580 and 𝜆7235 and Na iD /He i
become weaker, while the Fe ii 𝜆4924, 5018, 5169 lines become
stronger. During this period, the spectra do not show signs of O ii
lines. After 78 days, the Na iD /He i line vanishes while the flux in
the bluer part of the spectra shows a significant decrease, mainly due
to the line blanketing. After this epoch, a ’W’-shape profile is visible
around 4800 Å. This feature has been previously detected in several
SNe Ic (e.g. SN 2004aw; Taubenberger et al. 2006; SN 2007gr;
Valenti et al. 2008b; Hunter et al. 2009; Chen et al. 2014) and
SLSNe (e.g. SN 2015bn; Nicholl et al. 2016a).
To support our preliminary line identification in SN 2020wnt,
we employ the synow code (Fisher 2000) and the best quality spectra
at 49 days (−28 days from the 𝑟−band maximum) and 87 days
11 https://github.com/mnicholl/superbol/
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 7
4000 5000 6000 7000 8000 9000 10000
Rest Wavelength [A˚]
S
ca
le
d
F
lu
x
+
C
on
st
an
t
32d (-46d)
38d (-39d)
40d (-37d)
46d (-32d)
49d (-28d)
58d (-20d)
61d (-17d)
62d (-15d)
64d (-14d)
66d (-11d)
78d (+0d)
79d (+1d)
87d (+9d)
90d (+13d)
91d (+14d)
94d (+17d)
98d (+21d)
103d (+26d)
111d (+33d)
117d (+40d)
123d (+45d)
145d (+67d)
151d (+73d)
171d (+93d)
279d (+201d)
293d (+216d)
KAST AFOSC LRS ALFOSC YFOSC BFOSC OSIRIS
Ca II
Fe II
Fe II/Mg II
Fe II
Sc II
Sc II/Fe II
He I/Na I
Ba II
Si II
C II
C II O I Ca II
C I
[O III]
Mg I]
Hβ
[O III]
Fe II
[O I]
Na I
[O I]
Fe II
Hα
He I
[Ca II]+[O II]
Na ID (MW) Hα ⊕ ⊕[O III]
Figure 4. Spectral sequence of SN 2020wnt from 32 to 293 days from explosion in the rest-frame. The phases are labelled on the right. The numbers in
parentheses are the phases with respect to the maximum light in the 𝑟−band. Each spectrum has been corrected for Milky Way (MW) reddening and shifted
vertically by an arbitrary amount for presentation. The colour of the spectra represents the different instruments used to obtain the data. The vertical dashed lines
indicate the rest position of the strongest lines, vertical blue lines indicate emission lines from the host galaxy ([O iii] and H𝛼) and the narrow Na iD interstellar
feature from the MW, and the pink lines indicates the position of the telluric absorption (⊕ symbol). In the last two spectra, narrow emission lines from the
host galaxy (H𝛼, H𝛽, [O iii], [N ii], and [S ii], visible in grey) were removed for presentation.
(+9 days from the 𝑟−band maximum). For the spectrum before
peak, we assume a blackbody temperature of 𝑇bb = 10800 K and a
photospheric velocity of 𝑣ph = 8000 km s−1, while for the spectrum
after peak, we use a 𝑇bb = 8000 K and a photospheric velocity of
𝑣ph = 6000 km s−1. To reproduce the observed features in both
spectra, we include the lines of Ca ii, O i, Si ii, C ii, Na i, Fe ii, Sc ii,
Ba ii, Mg ii and Ti ii. As shown in Figure 5, the synthetic spectra at
49 and 87 days reproduce relatively well the observed features of
SN 2020wnt, allowing us to confirm the presence of C ii and Si ii.
Returning to the spectroscopic evolution, we see that from day
87 (+9 days from maximum), the Ca ii NIR triplet and O i become
stronger, and the region below 5500 Å is almost entirely dominated
MNRAS 000, 1–24 (2022)
8 Gutiérrez et al.
4000 5000 6000 7000 8000 9000
Rest Wavelength [A˚]
S
ca
le
d
flu
x
+
C
on
st
an
t
Ti II
Ca II
Fe II
Mg II +
Fe II
C II
Fe II
Na I
Sc II
Ba II
Si II
C II
C II O I Ca II
O I
SN 2020wnt at 49d (-28d)
SN 2020wnt at 87d (+9d)
SYNOW fits at 49 and 87d
Figure 5. Spectral comparison of SN 2020wnt at 49 and 87 days from
explosion (−28 and +9 from maximum, respectively) and the synow fits.
The synow synthetic spectra (green) are overplotted on the observed spectra
(black).
4500 5000
Rest Wavelength [A˚]
S
ca
le
d
fl
u
x
+
C
on
st
an
t
[O III]
Mg I]
Fe II
Hβ
[O III]
He I
SN 2020wnt at 279d
SN 2020wnt at 293d
6000 6500
Rest Wavelength [A˚]
Fe II
Hα
[O I]
Fe II
[N II]
Figure 6. Observed emission profiles of SN 2020wnt at 279 and 293 days
from explosion. Vertical lines indicate the rest position of the strongest lines.
Left panel: Emission profiles between 4200 and 5200 Å. Right panel:
Emission profiles between 6000 and 6800 Å.
by the iron-group lines. At 103 days, the C ii 𝜆6580 and 𝜆7235 lines
disappear while the continuum becomes redder. We detect a feature
at ∼ 9000 Å that is possibly due to C i 𝜆9183. As the temperature
decreases, C i lines start to be detected.
After 123 days (+45 from the peak), SN 2020wnt starts the
transition to the nebular phase. This is indicated by the Ca ii NIR
triplet, which shows signs of an emission component. From 123
to 171 days, the spectral evolution is slow. The most significant
difference is the strengthening of the emission component in the
Ca ii NIR triplet, as well as in the lines detected at ∼ 4600 Å,
∼ 6000 Å, and ∼ 7300 Å.
The last two spectra of SN 2020wnt were obtained at 279 and
293 days (+201 and +216 days from the peak), once the SN returned
from behind the Sun. At 279 days, the spectrum shows emission
lines of [Ca ii] 𝜆𝜆7291, 7324 with the possible contribution of [O ii]
𝜆𝜆7320, 7330), [O i] 𝜆𝜆6300, 6364, [O i] 𝜆5577, Na i, Mg i] 𝜆4571
plus Fe ii, a weak feature near ∼ 7100 Å, possibly caused by He i
𝜆7065, and a broad emission at ∼ 5000, which could be identified
as either the broad [O iii] 𝜆4959, 𝜆5007 components, Fe ii, He i or
1.0 1.5 2.0 2.5
Rest Wavelength [µm]
0.5
1.0
1.5
2.0
2.5
L
og
(F
λ
)
+
C
on
st
an
t
SYNOW fit
SN 2020wnt at 49d (-29d)
O I
Ca II
C I/O I/Mg II
C I/Mg II
C I/Ca II
⊕ ⊕
Figure 7. NIR spectrum of SN 2020wnt (black) obtained at ∼ 49 days from
explosion (−29 days from peak) with SpeX compared with the synow fit.
Regions of strong telluric absorption are masked in blue.
[Fe ii]. We explored all these identification scenarios and concluded
that this feature is probably caused by [O iii] 𝜆4959, 𝜆5007 with
some contribution of He i (see Figure 6). Despite the bluer part of
the spectrum is a bit noisy, we clearly see an emission line produced
by Ca ii H&K. An excess around 4300 Å is noticed and could be
consistent with [O iii] 𝜆4363.
One remarkable characteristic in this phase is the identification
of an emission line on the red side of [O i] 𝜆𝜆6300, 6364, which
could be caused either by H𝛼, Fe ii 𝜆6456, 6518 or [N ii] 𝜆𝜆6548,
6583. Fitting aGaussian, we find that the line is centred at∼ 6480Å.
If this feature is caused entirely by either H𝛼 or [N ii], a significant
blueshift (> 3000 km s−1) has to be taken into account, while for
Fe ii, the wavelength match is satisfactory. The possible detection
of Fe ii 𝜆6155 could support this hypothesis. However, a broad line
at the position of H𝛽 is also found, which suggests that H𝛼 should
contribute to the boxy feature. Therefore, this boxy emission line is
probably a result of blended lines of Fe ii, H𝛼, and probably [N ii].
A zoom around the H𝛼 and H𝛽 regions is presented in Figure 6.
The last spectrum, at 293 days, displays the same lines as
the spectrum at 279 days (from 5000 Å) plus emission features of
the Ca ii NIR triplet and O i 𝜆7774 (detected thanks to the larger
wavelength coverage). With the identification of these features, we
deduce that the strongest line at this phase is the Ca ii NIR triplet.
From 279 to 293 days, the spectra display twomajor changes, related
to the decrease in the intensity of the lines near 5300, 6155 Å (Fe ii)
and 6500 Å (the boxy emission line). The consistent decrease in the
strength of these lines supports the idea that the Fe ii 𝜆6456, 6518
lines should contribute to the line emission observed on the red side
of [O i] 𝜆𝜆6300, 6364.
Figure 7 shows a NIR spectrum obtained at ∼ 49 days from
explosion (−29 days from peak), covering a wavelength range be-
tween 0.63 and 2.5 𝜇m. At this phase, the spectrum displays a blue
continuum with clear features of O i, Ca ii, previously identified in
the optical spectra, plus multiple features of C i. We also detect an
absorption line at ∼ 1.05 𝜇m that could be either He i 1.083 𝜇m,
C i 1.069 𝜇m, or even a mix of them. If this line is due to He i, we
would also see a clear line detection near 2.058 𝜇m; however, this
is not the case. The lack of this line suggests that the contribution
of He i to the 1.05 𝜇m absorption is small. Similarly, the absorption
features at 1.26 and 2.12 𝜇m could be also produced by C i. There
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 9
0 50 100 150
Rest-frame days from explosion
2500
5000
7500
10000
12500
E
xp
an
si
on
ve
lo
ci
ty
[k
m
s−
1
]
Mg II λ4481
Si II λ6355
C II λ6580
O I λ7774
Ca II NIR T
−50 0 50
Rest-frame days from r−band maximum light
Figure 8. Evolution of expansion velocities derived from the minimum flux
of several absorption lines of SN 2020wnt (see Gutiérrez et al. 2017, for de-
tails in the estimation of velocities and their uncertainties). For comparison,
we also include the Si ii and O i velocities of the SLSN 2015bn (solid lines),
and the velocities of the carbon-rich type Ic SN 2007gr (dashed lines). For
this object, we stretched the time by a factor of 4 (𝑡 = 4 × 𝑡0). This factor is
obtained from the light curve analysis (see Section 4.3).
are several carbon lines near those wavelengths (see Millard et al.
1999; Valenti et al. 2008b; Hunter et al. 2009).
To identify the ions that produce the lines in the NIR, we create
a synthetic spectrum with synow by using the same parameters as
mentioned before. We consider a blackbody temperature of 𝑇bb =
10800 K and a photospheric velocity of 𝑣ph = 8000 km s−1. We
reproduce the observed spectrum only including Ca ii, O i, C i and
Mg ii, as shown in Figure 7, but beyond 1.8 𝜇m none of the ions
seems to produce absorption lines. The identification of C lines
in the NIR reinforces the idea that SN 2020wnt has a carbon-rich
progenitor.
3.5 Expansion velocities
We measure the expansion velocities of the ejecta for five spectral
lines (Mg ii, Si ii, C ii, O i and Ca ii NIR triplet) using the spectra
covering the phases from 32 and 151 days from explosion. These
velocities were obtained from the minimum flux of the absorption
component and are presented in Figure 8. In this analysis, we do not
include the iron lines (Fe ii 𝜆4924, 5018, 5169 Å) due to complica-
tions in estimating their velocities as these lines seem to be blended
with other ions. Figure 8 shows that the evolution of the expansion
velocities can be split into two groups: one characterised by an ini-
tial decrease that then flattens, and a second group that shows a
monotonic decline with time. In the first group, we find Ca ii NIR
triplet and O i, whereas Mg ii, Si ii and C ii are found in the second
group. In addition to this behaviour, we also see that the first group
shows higher velocities, suggesting that the Ca iiNIR triplet and O i
lines mostly form in the outer part of the ejecta, while the Mg ii,
Si ii and C ii lines form in the inner layers. The Ca ii NIR triplet has
the highest expansion velocities, decreasing from ∼ 12800 km s−1
at 32 days to ∼ 10100 km s−1 at 151 days. On the other hand, Si ii
decreases from 9600 km s−1 to just 2100 km s−1.
For comparison, in Figure 8we include the expansion velocities
of the carbon-rich type Ic SN 2007gr (Hunter et al. 2009) corrected
by a temporal factor of ∼ 4 in order to match the overall evolution
of SN 2020wnt (See Section 4.3), and the Si ii and O i velocities of
the SLSN 2015bn (Nicholl et al. 2016a). Overall, the velocities of
SN 2020wnt and SN 2007gr are comparable, except for that inferred
from the Ca ii NIR triplet, which has higher values in SN 2007gr.
When comparing SN 2020wnt with SN 2015bn, we see that the O i
velocities are very similar, but the Si ii velocities show a completely
different behaviour. While the velocities of SN 2015bn shows a
slow drop, in SN 2020wnt we notice a more rapid decrease. The
velocity values for all objects are low compared to those observed in
typical SNe Ibc (e.g. Prentice et al. 2019) and SLSNe, respectively.
Similarly, low velocities were also found for SN 2007gr (Hunter
et al. 2009) and SN 2015bn (Nicholl et al. 2016a).
4 COMPARISONWITH OTHER SNE
SN 2020wnt is a hybrid object sharing the properties of both SLSN
and SN Ic classes. The light-curve morphology is comparable to
that observed in several SLSNe (pre-peak bumps, long rise to the
main peak, slow evolution); with absolute magnitude at peak within
the luminosity distribution of SLSNe-I, although at the lower end
(De Cia et al. 2018; Angus et al. 2019). Despite this, SN 2020wnt
is brighter (∼ −20.5 mag) compared to standard SNe Ic (peak
absolutemagnitudes ranging between−17 and−18mag; e.g. Taddia
et al. 2018b), but lies in the luminosity range studied by Gomez
et al. (2022) for luminous SNe. On the other hand, the spectra are
more similar to the type Ic class than to SLSNe. This similarity
is founded on the absence of the O ii lines (one of the features
characterising SLSNe), and the strength of different lines, such as
the Ca iiNIR triplet, Si ii and C ii lines. A major difference between
SN2020wnt andSNe Ic is the slowevolution observed in the spectra,
which is consistent with SLSNe-I. Given these hybrid properties,
we compare SN 2020wnt with both SLSNe and SNe Ic. These are
well observed slow-evolving SLSNe, SLSNe with pre-peak bumps,
SLSNe with C ii lines in their spectra, and two very well sampled
type Ic SNe. Details of the comparison sample are presented in
Table 2.
4.1 Light curve comparison
Figure 9 shows the 𝑟−band absolute light curve of SN 2020wnt com-
pared to five well-sampled slow-evolving SLSNe-I and two normal
SNe Ic. From the comparison, the evolution of SN 2020wnt is sim-
ilar to that observed in SN 2017gci (this object has characteristics
of both slow- and fast-evolving SLSNe-I; Fiore et al. 2021). Ex-
cluding the type Ic SNe 2004aw and 2007gr, which are evidently
faint, SN 2020wnt is the faintest object in the (SLSN-I) sample. It
is > 1 mag fainter than the brightest object, SN 2015bn. Inspecting
the evolution at the early phases, we notice incomplete information
for the SLSN-I sample. Therefore, it is hard to know the rise-time
duration and how these objects evolve at very early phases, i.e., if
they show an initial peak (bump) or not. The only objects with an
early detection were PTF12dam and SN 2015bn. For SN 2015bn,
Nicholl et al. (2016a) estimated a rise time of 79 days, while for
SN 2020wnt, we estimate a rise time of ∼ 77.5 days in 𝑟. These rise
times are among the longest presented to date.
Analysing the shape of the light curve, we see that SN 2020wnt
and SN 2017gci have a similar behaviour. After maximum, the
decline slope of both objects changes at 50 − 55 days from peak.
Later, a shoulder is observed. For SN 2020wnt, after ∼ +50 days
from peak, the decline rate is comparable to that expected from
the 56Co decay (see Section 3.2). This evolution is observed until
MNRAS 000, 1–24 (2022)
10 Gutiérrez et al.
Table 2. Detailed properties of the comparison sample.
SN Redshift E(B−V)𝑀𝑊 Characteristics★ References†
(mag)
SLSNe
SN 2006oz 0.376 0.042 Bumpy (1)
SN 2007bi 0.127 0.028 Slow (2), (3)
PTF12dam 0.107 0.012 Slow; C ii (4), (5), (6)
LSQ14an 0.1637 0.074 Slow (7)
LSQ14bdq 0.345 0.056 Bumpy (8)
DES14X3taz 0.608 0.022 Bumpy (9)
SN 2015bn 0.1 0.022 Slow (10), (11)
DES15S2nr 0.22 0.030 Bumpy (12)
SN 2017gci 0.0873 0.116 Slow; C ii (13)
DES17X1amf 0.92 0.022 Bumpy (12)
SN 2018bsz 0.0267 0.214 Bumpy; C ii (14), (15), (16)
SNe Ic
SN 2004aw 0.0163 0.022 – (17)
SN 2007gr 0.0017 0.055 C ii (18), (19), (20)
★Characteristics: Slow: Slow-evolving SLSNe; Bumpy: SLSNe with pre-
peak bumps; C ii: C ii lines in the spectra.
†References: (1) Leloudas et al. (2012); (2) Gal-Yam et al. (2009); (3) Young
et al. (2010); (4) Nicholl et al. (2013); (5) Chen et al. (2015); (6) Vreeswijk
et al. (2017); (7) Inserra et al. (2017); (8) (Nicholl et al. 2015b); (9) Smith
et al. (2016); (10) Nicholl et al. (2016a); (11) Nicholl et al. (2016b); (12)
Angus et al. (2019); (13) Fiore et al. (2021); (14) Anderson et al. (2018);
(15) Chen et al. (2021); (16) Pursiainen et al. (2022); (17) Taubenberger
et al. (2006); (18) Valenti et al. (2008b); (19) Hunter et al. (2009); (20) Chen
et al. (2014).
−100 0 100 200 300 400
Days from Rest-frame maximum
−22
−21
−20
−19
−18
−17
−16
−15
M
r
(r
es
t-
fr
am
e)
56
Co decay
SN 2004aw
SN 2007bi
SN 2007gr
PFT12dam
LSQ14an
SN 2015bn
SN 2017gci
SN 2020wnt
Figure 9. Comparison of the 𝑟−band light curves of SN 2020wnt with the
slow-evolving SLSNe SN 2007bi, PTF12dam, LSQ14an SN 2015bn, and
SN 2017gci, and the type Ic SN 2004aw and SN 2007gr (given as reference).
All SLSNe have been K-corrected to rest-frame. Only corrections for Milky
Way extinction have been applied.
∼ +200 days from peak. After this, the slope changes again, showing
a very fast linear decline.
As mentioned in Section 3.2, the 𝑔𝑐𝑟−band light curves of
SN 2020wnt have an initial bump with a relatively short duration.
To analyse the pre-peak bump in SN 2020wnt, we compare the
𝑔−band early light curve with a sample of SLSNe that show signs
of an early bump. This comparison is presented in Figure 10. To
easily examine these objects, we arranged the light curves in terms
of the rise-time to the main peak. SN 2020wnt has the longest rise
time (∼ 72 days), while SN 2018bsz has the shortest one (around
−80 −60 −40 −20 0
Rest frame days from g−band peak
−20
−15
−10
M
g
(r
es
t
fr
am
e)
SN 2020wnt
LSQ14bdq + 4
DES15S2nr + 3.5
DES14X3taz + 6
SN 2006oz + 7
DES17X1amf + 9
SN 2018bsz + 9
Figure 10. Rest-frame, 𝑔−band light curve of SN 2020wnt compared with
SLSNe showing pre-peak bumps: SN 2006oz, LSQ14bdq, DES14X3taz,
DES15S2nr, DES17X1amf, and SN 2018bsz.
10 days). An opposite behaviour is observed in the duration of
the initial peak. Here, SN 2020wnt shows the shortest bump with
a duration < 5 days, while the longest initial bump is observed
for SN 2018bsz (> 25 days; see Anderson et al. 2018 for better
constraints in other bands). In terms of luminosity, we find that
LSQ14bdq is the most luminous object, DES15S2nr is the least
luminous, followed by SN 2020wnt and SN 2018bsz, which have
similar absolute magnitudes at peak.
4.2 Spectral comparison
The spectral comparison of SN 2020wnt and slow-evolving SLSNe-
I is presented in Figure 11. To examine the similarities and dif-
ferences between these objects, we select two reference epochs:
around peak (top) and ∼ +50 frommaximum light (bottom). Firstly,
we notice a large diversity. Around peak, SN 2020wnt has a dis-
tinct spectrum, dominated by strong absorption lines. In particular,
the W-shape profile due to Fe lines is a remarkable characteristic.
Though the spectral coverage does not allow to see the full profile
of the Ca ii NIR triplet, based on the spectra before and after peak
(Figure 4), this line is intense in SN 2020wnt, but it seems to be
absent in the other objects. Only a few common features are identi-
fied among SN 2020wnt and the comparison sample: C ii lines with
SN 2017gci, and Si ii and O i lines with SN 2015bn.
At around +50 days from maximum light, the spectra are still
quite heterogeneous, although SN 2020wnt and SN 2017gci look
more alike than before. This is seen in the bluer part, with the
blended feature due to the Mg ii and Fe ii lines, and in the redder
part, with the detection of strong features of O i and Ca ii NIR
triplet. Both SN 2020wnt and SN 2015bn still share similar profiles
of Si ii and O i. Fe ii and O i are visible in all objects. In contrast
to the photometric behaviour, the spectrum of SN 2020wnt seems
to evolve slower than that of SN 2007bi and LSQ14an. At the later
phase, these objects show clear signs of nebular lines (e.g. [Ca ii]
+ [O ii]), while SN 2020wnt still shows lines of the photospheric
phase.
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 11
L
og
(F
λ
)
+
C
on
st
an
t
Ca II
Fe II
Na ID
Si II
C II
C II
O I Ca II
SN 2020wnt
(0d)
SN 2017gci
(-3d)
SN 2015bn
(-2d)
PTF12dam
(-2d)
4000 6000 8000 10000
Rest Wavelength [A˚]
L
og
(F
λ
)
+
C
on
st
an
t
Ca II
Fe II+
Mg II
Fe II
Na ID
Si II O I Ca II
C I
[Ca II]
SN 2020wnt
(+45d)
SN 2017gci
(+51d)
SN 2015bn
(+50d)
LSQ14an
(+50d)
PTF12dam
(+55d)
SN 2007bi
(+56d)
Figure 11. SN 2020wnt at around maximum light (top) and at ∼ +50
days from 𝑟−band maximum (bottom) compared with well-sampled slow-
evolving SLSNe: SN 2007bi, PTF12dam, LSQ14an, SN 2015bn, and
SN 2017gci. Each spectrum has been corrected for Milky Way redden-
ing and shifted vertically by an arbitrary amount for presentation. Vertical
lines indicate the rest position of the strongest lines. The phases and SN
names are labeled on the right.
4.3 Comparison with SN 2015bn and SN 2007gr
We now compare SN 2020wnt with the most extreme objects shown
in Figure 9: SN 2015bn, the brightest object in the comparison
sample and one of the best observed SLSNe to date, and the carbon-
rich type Ic SN 2007gr, the faintest SN in the plot, and the best-
spectral match object found by gelato (Harutyunyan et al. 2008).
The spectral comparison between SN 2020wnt and SN 2015bn
at seven different epochs is presented in Figure 12. Before the peak
(top panel), the spectra of both objects are characterised by a blue
continuum with lines of O i, Si ii, C ii and Fe ii. In SN 2020wnt,
these features are stronger at all phases. In contrast, the main differ-
ences lie in the O ii and Ca ii lines. The spectra of SN 2020wnt do
not show the W-shape O ii features as observed in SN 2015bn, and
other SLSN-I events. The presence/absence of these lines depends
on the temperature (more details in Section 6). Unlike SN 2015bn,
SN 2020wnt shows very strong Ca ii absorptions. These Ca ii ab-
sorption features are not observed in young SLSNe-I. After peak
(bottom panel), in contrast with SN 2015bn, SN 2020wnt has a red-
der continuum, more intense lines and the bluer part of the spectrum
is dominated by Fe-group lines. Some signs of emission components
are also detected, suggesting the beginning of the transition to the
nebular phase. All these properties are delayed in SN 2015bn.
In Figure 13, we compare SN 2020wnt and the type Ic carbon-
rich SN 2007gr. Although the light curves (top panel) of these two
objects are completely different at the first glance, they share a
similar spectroscopic evolution (bottom panel) with some time lag
L
og
(F
λ
)
+
C
on
st
an
t SN 2020wntSN 2015bn
Ca II Fe II
Na I
Si II
C II
C II
O I
Ca II
-29d
-27d
-20d
-20d
-12d
-9d
O II
4000 6000 8000 10000
Rest Wavelength [A˚]
L
og
(F
λ
)
+
C
on
st
an
t +1d
+3d
+17d
+20d
+70d
+72d
+96d
+106d
Figure 12. Spectral comparison of SN 2020wnt with the slow-evolving
SN 2015bn at eight different epochs before (top panel) and after (bottom
panel) the maximum light in the 𝑟−band. Each spectrum has been corrected
for Milky Way reddening and shifted vertically by an arbitrary amount for
presentation. Vertical lines indicate the rest position of the strongest lines.
The phases are labeled on the right.
(𝑡 = 4 × 𝑡0; where 𝑡0 is the rest-frame maximum of SN 2007gr),
which was initially found by the spectral matching from gelato.
From the light curves in the 𝑟−band, we estimate that SN 2020wnt
is ∼ 3 mag brighter than SN 2007gr around the maximum light.
SN 2020wnt also has a pre-peak bump light curve morphology and
evolves on a much longer timescale than its fainter counterpart. In
order to have a broad light curve similar to that of SN 2020wnt, we
applied a temporal correction to SN 2007gr. We find that a factor
of 4 can reproduce the light curve width of SN 2020wnt, and it is,
in turn, in agreement with that found by the spectral matching. This
correction is included in the top panel of Figure 13 (open orange
circles).
In the bottom panels of Figure 13, the spectroscopic com-
parison between SN 2020wnt and SN 2007gr is presented at four
different epochs. Analysing the spectral features, one sees that the
main similarity is the detection of the carbon lines, while the main
difference is the strength of the Na i line, which is more intense
in SN 2007gr. From the spectral matching with gelato, we found
a time lag of 𝑡 = 4 × 𝑡0. Therefore, the SN 2020wnt spectrum at
−27 days from peak is compatible with that of SN 2007gr at −7
days from peak. Both SNe show strong features of Ca ii, O i, and
the clear signs of Si ii, C ii and the W-shape Fe ii lines at around
4800 – 5200Å. After peak, both objects evolve maintaining the time
lag identified in the spectra before peak, which is consistent with
that obtained from the light curves. The similarity in the spectral
evolution suggests that both objects may arise from a carbon-rich
MNRAS 000, 1–24 (2022)
12 Gutiérrez et al.
−100 0 100 200 300 400
Days from Rest-frame maximum
−20
−18
−16
−14
M
r
(r
es
t-
fr
am
e) 56Co decay
56Co decay
SN 2020wnt
SN 2007gr
SN 2007gr (t = 4× t0; Mr=Mr0−2.5)
4000 6000 8000 10000
Rest Wavelength [A˚]
L
og
(F
λ
)
+
C
on
st
an
t
SN 2020wnt
SN 2007gr
Ca II
Fe II Na ID/C II
Si II
C II
C II
O I Ca II
C I
-28d
-7d (−28d)
+26d
+7d (+28d)
+73d
+20d (+80d)
+93d
+29d (+116d)
Figure 13. Light curve and spectral comparison with the type Ic carbon-rich
SN 2007gr.Top:Comparison of the 𝑟−band light curves. For the light curve
of SN 2007gr, we applied a temporal correction (𝑡 = 4× 𝑡0) to reproduce the
same shape and width as SN 2020wnt. A shift of 2.5 mag is applied to the
magnitude in order to have a similar brightness in the transition and the tail
(from∼ +53 days from peak).Bottom: Spectral comparison at four different
epochs from the 𝑟−band maximum. Each spectrum has been corrected for
Milky Way reddening and shifted vertically by an arbitrary amount for
presentation. Vertical lines indicate the rest position of the strongest lines.
The rest-frame phases are labeled on the right, while the epochs with the
temporal correction (for SN 2007gr, 𝑡 = 4 × 𝑡0) are in parenthesis. This
correction gives a consistent epochs to SN 2020wnt.
progenitors, however, the longer timescale and the brightness of the
light curve may suggest different explosion parameters.
In Figure 14, we compare the 𝑔 − 𝑟 and 𝑟 − 𝑖 colours of
SN 2020wnt, SN 2015bn and SN 2007gr (applying the temporal
correction previously mentioned). In 𝑔 − 𝑟, we find that from the
maximum light, the evolution of SN 2020wnt and SN 2015bn are
very different. While SN 2020wnt evolves quickly to the red, reach-
ing a peak ∼ 50 days later, SN 2015bn remains bluer a much longer
time. At later phases, the colours are more alike. On the other hand,
the 𝑟 − 𝑖 colours are almost identical in the three objects up to +150
days from the maximum (SN 2007gr being the bluest). After this,
the evolution diverges. For instance, SN 2020wnt and SN 2007gr
become a bit bluer, whereas SN 2015bn evolves to the red. In Fig-
ure 14, we also include the intrinsic colour templates of SNe Ic
from Stritzinger et al. (2018). To these templates, we apply the
same temporal corrections as that found for SN 2007gr. From this
−0.5
0.0
0.5
1.0
1.5
g
−
r
SN 2020wnt
Template (SNe Ic)
SN 2015bn
−50 0 50 100 150 200 250
Rest-frame days from r−band maximum light
0 50 100 150 200 250 300
Rest-frame days from explosion
−0.5
0.0
0.5
1.0
r
−
i
SN 2020wnt
Template (SNe Ic)
SN 2007gr
SN 2015bn
Figure 14. 𝑔 − 𝑟 and 𝑟 − 𝑖 colour curves of SN 2020wnt. For comparison,
in solid lines (purple and green) we show the intrinsic SNe Ic colour-curves
templates (𝑔 − 𝑟 and 𝑟 − 𝑖) from Stritzinger et al. (2018) multiplied by
a temporal factor of 4 (𝑡 = 4 × 𝑡0). We also include the 𝑟 − 𝑖 colour of
SN 2007gr (multiplied by the temporal factor; dashed orange line) and the
𝑔 − 𝑟 and 𝑟 − 𝑖 colours of SN 2015bn (dashed light red line; Nicholl et al.
(2016a)). The data used here for SN 2007gr were taken from Bianco et al.
(2014).
comparison, 1) using these templates to constrain the host extinc-
tion, we find that for SN 2020wnt it is negligible; 2) the colour
of SN 2020wnt between the maximum light and +40 days from
the maximum (temporal correction included), evolves similarly as
SNe Ic in both 𝑔 − 𝑟 and 𝑟 − 𝑖. This suggests that the colour of
SN 2020wnt is similar to that observed in standard SNe Ic but on a
longer timescale.
In Figure 15, we compare the late phase spectra of SN 2020wnt
with those of SN 2007gr (top) and SN 2015bn (bottom) at similar
epochs. As before, we find that the comparison with SN 2007gr has
a time lag of 𝑡 = 4 × 𝑡0, i.e., the spectrum of SN 2020wnt at +201
days from peak is similar to the spectra of SN 2007gr after 50 days
from peak. In the top panel of this Figure, we compare our spectra
at +201 days (blue part, left panel) and +216 days (red part, right
panel) with SN 2007gr at +51, and +78 days from peak (with the
correction the epochs are +204, and +312). Analysing the blue part
(left panel), we see that both objects have the similar lines, but they
seem to be stronger in SN 2020wnt. The shoulder on the red side
of [O i] 𝜆𝜆6300, 6364 is also detected in SN 2007gr, though a bit
fainter. From the red part (spectrum at +216 days, right panel), the
lines observed in both objects fit very well. The larger discrepancy
is in the intensity of the Ca ii NIR triplet feature, which is stronger
at all epochs in SN 2007gr. At this phase, the shoulder of the red
side of the [O i] line matches better than before.
In the bottom panel of Figure 15, we present the late spectra
of SN 2020wnt and SN 2015bn. For SN 2015bn we use two epochs
(+243 and +315 days from peak) that have some similarities with
our spectra. Examining the blue part of the spectra, we notice some
differences between these two objects. For instance, in the three
spectra of SN 2015bn, there are no signs of a feature on the red
side of [O i] 𝜆𝜆6300, 6364, of the emission around 5300 Å (at-
tributed to Fe ii) and the line at ∼ 7100 Å (possibly He i) detected
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 13
S
ca
le
d
F
lu
x
+
C
on
st
an
t
Ca II
[O III]
Mg I]
Hβ
[O III]
Fe II
[O I]
Na I
[O I]
Fe II
Hα
He I
[Ca II]
SN 2020wnt at +201d
SN 2007gr at +51d (204d)
SN 2007gr at +78d (312d)
Fe II
[O I]
Na I
[O I]
Fe II
Hα
He I
[Ca II]
O I
Ca II NIRSN 2020wnt at +216d
SN 2007gr at +51d (204d)
SN 2007gr at +78d (312d)
4000 5000 6000 7000
Rest Wavelength [A˚]
S
ca
le
d
F
lu
x
+
C
on
st
an
t
Ca II
[O III]
Mg I]
Hβ
[O III]
Mg I
Fe II
[O I]
Na I
[O I]
Fe II
Hα
He I
[Ca II]
SN 2020wnt at +201d
SN 2015bn at +243d
SN 2015bn at +315d
5000 6000 7000 8000 9000 10000
Rest Wavelength [A˚]
Mg I
Fe II
[O I]
Na I
[O I]
Fe II
Hα
He I
[Ca II]
+ [O II]
O I
Ca II NIR
SN 2020wnt at +216 d
SN 2015bn at +243d
SN 2015bn at +315d
Figure 15. Spectral comparison between SN 2020wnt and SN 2007gr (top) and SN 2020wnt and SN 2015bn (bottom). Each spectrum has been corrected for
Milky Way reddening. Spectra are scaled to match the [Ca ii] line. Vertical lines indicate the rest wavelengths of the strongest lines. The phases are labeled in
the legend. In SN 2020wnt, the narrow emission lines from the host galaxy were removed for presentation. For SN 2007gr the phases in brackets correspond
to the epochs with the temporal corrections (𝑡 = 4 × 𝑡0).
in SN 2020wnt. Now, the red part (right panel) is similar, although
SN 2020wnt has a stronger Ca ii NIR triplet, but a weaker O i than
SN 2015bn at all epochs. Overall, the SN 2015bn spectrum that
best matches SN 2020wnt is that at +315 days. The similarity of
the nebular spectra of SN 2020wnt with those of SN 2007gr and
SN 2015bn suggests that these objects are related, and they possibly
have a common origin, including a similar chemical composition.
5 EXPLOSION AND PROGENITOR SCENARIOS
5.1 Nebular properties
To constrain the internal conditions and core structure of
SN 2020wnt, we investigate the nebular spectra in detail. Based
on the findings of Jerkstrand et al. (2014), the O i mass responsible
for the line emission can be estimated from the oxygen luminosity,
as follows:
M𝑂 =
𝐿6300,6364/𝛽6300,6364
9.7×1041 erg s−1 × 𝑒𝑥𝑝( 22720 𝐾𝑇 ) M
where L6300,6364 is the line luminosity of [O i] 𝜆𝜆6300, 6364,
𝛽6300,6364 is the is the Sobolev escape probability, and T is the
temperature. Here, the temperature can be derived from the [O i]
𝜆5577 to [O i] 𝜆𝜆6300, 6364 ratio (Houck & Fransson 1996; Jerk-
strand et al. 2014). Given that the SN 2020wnt spectra are not fully
nebular, the estimation of these parameters depends on how we de-
fine the line fluxes, and in turn, the temperature. To estimate the
luminosities, a linear fit to the continuum is subtracted, and then,
we fit a Gaussian to the line. As [O i] 𝜆5577 is hard to measure,
we take the extreme values (minimum and maximum flux) obtained
from the Gaussian fit. Thus, using the spectrum at +216 days from
the maximum light, and assuming a 𝛽5577/𝛽6300,6364 ratio equal
to 1, we get a temperature between ∼ 3800 and ∼ 4500 K. For
𝛽6300,6364 = 1 (i.e. assuming optically thin emission), we can ob-
tain the minimum mass of oxygen required to produce the observed
[O i]. With these values, we derive a M𝑂 ≈ 2 − 4 M . Using a
similar approach, Nicholl et al. (2016b) found M𝑂 ≈ 9M for SN
2015bn, while Mazzali et al. (2010) found M𝑂 ≈ 1 M by mod-
elling the nebular spectra of SN 2007gr. Despite the uncertainties in
the estimation of the O imass, the values found for SN 2020wnt are
intermediate between those derived for SN 2015bn and SN 2007gr.
Information on the coremass is usually inferred from the [Ca ii]
𝜆𝜆7291, 7324/[O i]𝜆𝜆6300, 6364 ratio (Fransson&Chevalier 1989;
Elmhamdi et al. 2004; Kuncarayakti et al. 2015), although we are
aware that the ratio is sensitive to various parameters (e.g., Li &
McCray 1993; Jerkstrand 2017; Dessart et al. 2021). Nonetheless,
we calculate this ratio in order to compare it with objects from
the literature. Using the spectrum at +216 days, we compute a
[Ca ii]/[O i] flux ratio of ∼ 1.1, which is within the range of values
found for core-collapse SNe (lower than 1.43; e.g., Kuncarayakti
et al. 2015; Terreran et al. 2019; Gutiérrez et al. 2020b), but larger
than the value found for SN 2015bn (0.5; Nicholl et al. 2016b).
Although a [Ca ii]/[O i] ratio of ∼ 1.1, suggests a relatively low
helium core mass, we note that the spectrum at +216 days is not
completely nebular. Therefore, the ratio estimation can be affected
by this issue, as seen in some core-collapse SNe with good nebular
coverage (e.g. SN 2017ivv; Gutiérrez et al. 2020b). Additionally,
we also highlight that the [Ca ii] feature could be contaminated by
[O ii]. In fact, if the nebular feature around 7300 Å is dominated by
[O ii] instead of [Ca ii] (e.g. SN 2007bi; Gal-Yam et al. 2009), this
ratio cannot be a good proxy of the progenitor core mass. In this
case, the real O amount could be much higher, and more consistent
with the very high value of the 56Ni inferred in Section 5.2.
In Section 3.4, we pointed out that the strongest line observed
in SN 2020wnt at +216 days from peak (293 days from explosion)
was the Ca ii NIR triplet. From the nebular comparison, we found
this feature is fainter in SN 2020wnt than in SN 2007gr but stronger
than in SN 2015bn. By modelling SLSN nebular spectra, Jerkstrand
et al. (2017) found that an electron density of 𝑛𝑒 & 108 cm−3 is
needed to reproduce a Ca iiNIR/[Ca ii] larger than 1. For SN 2015bn
they derived a ratio of 1.7, which is unusually high for the nebular
phase. Following this approach, we measure the Ca ii NIR/[Ca ii]
ratio for SN 2020wnt and we find a value of 3.3. This value is
MNRAS 000, 1–24 (2022)
14 Gutiérrez et al.
twice that derived for SN 2015bn. This suggests extraordinarily
high electron densities for SN 2020wnt. However, as mentioned
before, the spectrum at +216 days is not fully nebular and the ratios
we measure might correspond to some limits.
In the spectrum of SN 2020wnt at +201 days, we recognise two
broad features around 4350 and 5000 Å. These features, that could
be attributed to [O iii] 𝜆4363 and [O iii] 𝜆𝜆4959, 5007, have been
detected in several core-collapse SNe at very late phases (e.g. Fesen
et al. 1999; Milisavljevic et al. 2012) and in a few SLSNe during the
nebular phase: PS1-14bj (Lunnan et al. 2016), LSQ14an (Inserra
et al. 2017), SN 2015bn and SN2010kd (Kumar et al. 2020). As seen
in Figures 4, 6 and 15, SN 2020wnt exhibits a relatively strong [O iii]
𝜆𝜆4959, 5007, but a weak [O iii] 𝜆4363, which suggest a high [O iii]
𝜆𝜆4959/ [O iii] 𝜆4363 flux ratio. This ratio can provide information
about the temperature and electron density of the emitting region
where they formed. Fitting a Gaussian to these lines, we measure a
[O iii] 𝜆𝜆4959, 5007/ [O iii] 𝜆4363 flux ratio of ∼ 3.5. This value is
much larger than those measured for PS1-14bj (Lunnan et al. 2016)
and LSQ14an (Inserra et al. 2017), and suggests electron densities
greater than 106 cm−3 (Fesen et al. 1999; Jerkstrand et al. 2017).
Summarising, from the line ratio analysis (e.g. Ca iiNIR/[Ca ii]
and [O iii] 𝜆𝜆4959/, 5007 [O iii] 𝜆4363), we derive high electron
densities (𝑛𝑒 ∼ 106 – 108 cm−3).Density values of around 108 cm−3
have been inferred before for the SNe IIn SN 1995N (Fransson et al.
2002) and SN 2010jl (Fransson et al. 2014), and more recently for
the SLSN 2015bn (Jerkstrand et al. 2017).
5.2 Light curve modelling
We explore several models, trying to explain the light curve mor-
phology of SN 2020wnt. Two important characteristics put some
constraints in our modelling: 1) the long-rise time to the main peak
and 2) the luminosity following the radioactive decay at 140 days
from explosion. These two properties point out that this SN may
belong to the rare class of SLSNe-I that are possibly powered by a
large amount of nickel production, similar to SN 2007bi (Gal-Yam
et al. 2009). Here, we also analyse the possibility that the main peak
can be powered by a magnetar.
To calculate the light curve and photospheric velocity, we use
a one-dimensional local thermodynamical equilibrium (LTE) radi-
ation hydrodynamical code presented in Bersten et al. (2011). The
code follows the complete evolution of the light curve in a self-
consistent way from the shock propagation to the nebular phase.
The explosion initiates by injecting some energy near the progen-
itor core. This produces a powerful shock wave that propagates
inside the progenitor until it arrives at the surface, where the pho-
tons begin to diffuse out. The code assumes a grey transport for
gamma photons produced during the radioactive decay, but allows
any distribution of this material inside the ejecta and assumes an
opacity of 𝜅𝛾 = 0.03 cm2 g−1. Therefore, we are able to calculate
the gamma-ray deposition in each part of the ejecta and estimate the
gamma-ray escape as a function of time. The inclusion of a mag-
netar as an extra source to power the SN, including the relativistic
effect, is accounted in the code (see Bersten et al. 2016, for details).
An extensive exploration of magnetar parameters for different types
of progenitors was discussed in Orellana et al. (2018).
An initial structure in hydrostatic equilibrium which simulates
the condition of the star at the pre-SN stage is needed to start the
hydrodynamical calculations. Here we assume both parametric and
stellar evolutionary models. The initial emission and the rise time
to the main peak of SN 2020wnt may indicate that a more extended
progenitor is required than the typical stripped (or compact) star
41.5
42.0
42.5
43.0
43.5
lo
g 1
0
L
bo
l
(e
rg
s−
1
)
SN 2020wnt
M25R50E7Ni3.5
M80R15E45Ni4 (κγ
reduced at t < 115d)
M80R15E45Ni4
(κγ constant)
0 20 40
Rest-frame days from explosion
42
43
lo
g
1
0
L
b
o
l
(e
rg
s−
1
)
0 100 200 300
Rest-frame days from explosion
5000
10000
V
el
oc
it
y
[k
m
s−
1
]
Mg II λ4481
Si II λ6355
C II λ6580
O I λ7774
Ca II NIR T
M25R50E7Ni3.5
M80R15E45Ni4
Figure 16. Top panel: Comparison between our best hydrodynamical mod-
els and the SN observations (grey stars). For the more massive model (ma-
genta), we use a progenitor with a pre-SN mass of 80 M , 𝐸exp = 45 ×1051
erg, a 56Nimass of 4M and an initial radius of 15R , while the lessmassive
model (green solid line) has a a pre-SN mass of 25 M and 𝐸exp = 7 ×1051
erg, a 56Ni mass of 3.5 M and an initial radius of 50 R . For the more
massive model, we show the light curve with 𝜅𝛾 constant (solid line), and a
light curve with 𝜅𝛾 reduction (dashed-line). The inset plot shows the light
curves at very early phases. Bottom panel: Photospheric velocity of the
models compared with the velocities of different species in SN 2020wnt.
assumed for H-free objects. Therefore, we test models typically
used for H-rich SNe with extended and dense envelopes like red
supergiant (RSG) or blue supergiant (BSG) stars, but manually
modifying the chemical abundance to produce a H-free envelope.
The RSG models are computed by stellar evolutionary calculation
while the BSGprogenitor is computed assuming a double polytropic
model given that it not easy to generate these BSG progenitors with
stellar evolution models.
Assuming only a radioactive source, we cannot find a suitable
solution for RSG (or stripped-envelope) progenitors, but models
with a structure similar to those used for 87A-like objects or BSG
progenitors (with our altered chemical composition) seem to offer
a good solution. We generate several initial configurations for dif-
ferent values of the pre-SN mass (between 15 and 100 M) and
radius (in the range of 15 to 80 R). We explore many values of the
explosion energy and 56Ni production for each configuration. Our
best models are shown in the top panel of Figure 16 and correspond
to two progenitors with a pre-SN radius of 50 and 15 R , pre-SN
masses of 25 M (green) and 80 M (magenta), and explosion en-
ergies of 7× 1051 erg and 45× 1051 erg, respectively. In both cases,
a large amount of 56Ni (3.5 and 4 M) is required to reproduce
the main peak of the light curve. For the more massive model (80
M), we slightly modified the gamma-ray opacity from 𝜅𝛾 = 0.03
to 𝜅𝛾 = 0.013 cm2 g−1 at around 115 days from the explosion to
improve the fit from this epoch (i.e. allowing an easier leakage of
gamma-ray photons; see Gutiérrez et al. 2021, for more details).
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 15
However, we also include this model but with a 𝜅𝛾 constant (solid
magenta line).
As observed in several SNe Ib/c (e.g. Sollerman et al. 2000,
and references therein), the gamma-ray leakage significantly affects
the light curve at late phases. More precisely, in several energetic
SNe (e.g. SNe Ic broad line), the nickel mass required to explain
the peak luminosity generally overestimates the tail luminosity (e.g.
Maeda et al. 2003). To partially solve this issue, it has been pro-
posed that an enhancement in the gamma-ray escape may happen
in latter epochs (after the main peak). This could be due to possible
asymmetries in the ejecta as the presence of low-density (or clumps)
zones (Tominaga et al. 2005; Folatelli et al. 2006). Such types of
structures could be produced by jets (e.g. Soker 2022) or Rayleigh-
Taylor instabilities. In these low-density regions, the gamma-rays
could escape more efficiently, which can be simulated by reducing
the kappa-gamma values. In addition, the grey transfer assumed for
the gamma-rays can require a time varying factor, as shown in Wilk
et al. (2019) by solving the relativistic radioactive transfer equation
for gamma-ray in SNe.
The photospheric velocities of the models are compared to the
velocities of different species in the bottom panel of Figure 16. The
less massive model (25 M) reproduces the observables exception-
ally well, particularly for Mg ii, Si ii and C ii; however, this is not
the case for the 80 M model. This model overestimates most of
the velocities but has a good agreement with O i from 100 to 160
days.
Regarding the light curve, we found that the early emission can
be attributed to the cooling phase for an extended progenitor of∼ 50
and 15R , respectively. However, none of themodels can reproduce
the luminosity at phases later than ∼ 245 days, since the hypotheses
used in the code fail at these late times. At these epochs, the lumi-
nosity declines slower than the 56Co decay (for about 30 days) and
then suddenly drops. Between ∼ 245 and ∼ 275 days, the excess
in luminosity is probably due to CSM interaction, which could be
supported by the presence of H emission lines (H𝛼 and H𝛽) in the
late spectra (see Sections 3.4). The sudden drop in luminosity may
be explained by the end of this interaction phase or alternative by
dust formation. When the SN reaches the minimum value in this de-
crease, it almost recovers the luminosity expected from radioactivity
(80M model). Ignoring these late epochs, the 25 and 80M mod-
els well describe the overall light curve evolution. Considering the
photospheric velocities, the 25 M model better represents the ob-
served properties of SN 2020wnt. However, given the large amount
of 56Ni (3.5M), this model turns out to be unrealistic. Relaxing the
velocity constraints and considering that reaching an 𝐸 ≥ 1052 re-
quires a much higher and possibly nonphysical neutrino deposition
fraction (Janka 2012; Terreran et al. 2017), we find that the 80 M
model reproduces the light curve properties of SN 2020wnt very
well. Additionally, given that t𝑛𝑒𝑏 ≈
√︁
𝜅 ∗ 𝑀/𝑣2 ≈ 𝑀 ∗ √︁𝜅/𝐸𝑘𝑖𝑛
and t𝑝𝑒𝑎𝑘 ≈
√︁
𝜅 ∗ 𝑀/(𝑣 ∗ 𝑐), we have t2
𝑛𝑒𝑏
≈ 𝑡2
𝑝𝑒𝑎𝑘
∗ 𝑐/𝑣. Since
t𝑝𝑒𝑎𝑘 is fixed and our more massive model has higher velocities
than the less massive model, the former is more compatible with
the expected values for t𝑛𝑒𝑏 . Therefore, SN 2020wnt is consistent
with a massive progenitor (pre-SN mass of 80 M), with a radius
of 15 R , explosion energy of 45 × 1051 erg, and 4 M of 56Ni.
These characteristics are compatible with the PISNe scenario (e.g.
Gal-Yam et al. 2009; Kozyreva et al. 2014, but see Dessart et al.
2013).
As mentioned above, we also analyse the possibility that
SN 2020wnt is powered by a magnetar. For this, we apply a ver-
sion of our one-dimensional LTE radiation hydrodynamics code
41.5
42.0
42.5
43.0
43.5
lo
g 1
0
L
bo
l
(e
rg
s−
1
)
SN 2020wnt
Model He8
Model He18
0 20 40
Rest-frame days from explosion
42
43
lo
g
1
0
L
b
o
l
(e
rg
s−
1
)
0 100 200 300
Rest-frame days from explosion
5000
10000
V
el
oc
it
y
[k
m
s−
1
]
Mg II λ4481
Si II λ6355
C II λ6580
O I λ7774
Ca II NIR T
Model He8
Model He18
Figure 17. Top panel: Comparison between the magnetar models and the
SN observations (grey stars). For the compact configuration (cyan solid
line), we use a progenitor with a pre-SN mass of 8 M , 𝐸exp = 1 × 1051
erg and a magnetar with an initial period 𝑃 ∼ 4.2 ms and a magnetic field
𝐵 ∼ 2×1014 G. In a more extended configuration (pink dashed-line), we use
a progenitor with a pre-SN mass of 18 M , 𝐸exp = 1 × 1051 erg, 𝑃 ∼ 2 ms
and a magnetic field 𝐵 ∼ 3 × 1014 G. The inset plot shows the light curves
at very early phases. Bottom panel: Photospheric velocity of the magnetar
models compared with the velocities of different species in SN 2020wnt.
that takes into account the power provided by a newborn magnetar
that loses rotational energy. That energy is fully deposited in the
inner zones of the ejecta. We use the ¤𝐸 vacuum prescription that
became popular after the work of Kasen & Bildsten (2010) with
the standard assumption of a magnetic dipole and a breaking index
𝑛 = 3. Modern and detailed treatments by Vurm & Metzger (2021)
indicate that this approach is roughly valid at early times from the
explosion. The parameters of this source are the surface magnetic
field strength 𝐵 and the initial rotation period 𝑃.
We test some of the H-poor progenitors, with a usual ra-
dioactive content of around 0.1 M . For the compact configura-
tion (𝑀 = 8 M with 𝑅 = 1.3 R and an explosion energy of
𝐸exp = 1 × 1051 erg ) the 𝐿peak of SN 2020wnt can be reached
with 𝑃 ∼ 4.2 ms and 𝐵 ∼ 2 × 1014 G, but the rise and post peak
are not well reproduced (see Figure 17). We explore the same con-
figuration with the mixing of the 56Ni up to a fraction of 0.95 of
the ejecta, but that makes no major changes. Then, we explore the
possibility of a magnetar being applied to a more extended RSG
structure. We use the results from Orellana et al. (2018) to select a
possible model for SN 2020wnt. In Figure 17, we present the results
for a progenitor with a pre-SN mass of 18 M and a radius of 725
R . In this case the H is converted to He artificially, to have an
inflated structure for a H-free star, though it is not the result of stel-
lar evolution calculations. In order to obtain the time scale and the
energy of the peak, the following magnetar parameters are required:
𝑃 ∼ 2 ms and 𝐵 ∼ 3 × 1014 G, assuming an explosion energy of
𝐸exp = 1 × 1051 erg. While this model provides an improvement
with respect to the compact model, it produces a worse represen-
MNRAS 000, 1–24 (2022)
16 Gutiérrez et al.
tation of the light curve than the 56Ni powered models presented
in Figure 16. Additionally, the velocities are not well reproduced.
Therefore, this result disfavours a magnetar model to explain the
overall light curve evolution of SN 2020wnt.
We remark that none of the light curve models presented here
is designed to fit the drastic drop in brightness observed from 273
days. Finally, although the 56Ni powered models presented in Fig-
ure 16 provide a good explanation of SN 2020wnt for both the
light curve and the velocities, the progenitor models used were built
parametrically (both the pre-SN density profile and the chemical
composition). It would be desirable to investigate what kind of evo-
lutionary path, if any, can generate this type of structure to provide
a more solid physical framework for our modelling.
5.3 CSM interaction
Despite the spectra of SN 2020wnt show remarkable similarities
with SN 2007gr and non-interacting SNe Ic, the light curve fluctu-
ations and the very high luminosity can be comfortably explained
with the powering of ejecta-CSM interaction. Under this premise,
we could consider that the actual 56Ni-powered SN light curve
would be below the observed one (similar to that observed for sev-
eral SNe Ic; see top panel of Figure 13) dominated by ejecta-CSM
interaction. We use TigerFit12 (Chatzopoulos et al. 2016; Wheeler
et al. 2017) to explore this scenario by assuming a steady-state wind
and a constant-density CSM shell, respectively. Although none of
the fits reproduces the observed light curve of SN 2020wnt, we
found that a steady-state wind CSM can reach a better approxima-
tion. In this context, a potentially pervasive interaction phase lasts
for ∼ 273 days and ends when the luminosity suddenly drops. Here,
the CSM is probably H and He free. After the drop in brightness,
the light curves continue to experience small fluctuations (see Fig-
ure 2), most likely attributed to episodes of CSM interaction. We
note that very massive stars (M𝑍 𝐴𝑀𝑆 ∼ 100−140M; Heger et al.
2003) are expected to lose mass through pulsational pair-instability.
The SN ejecta are then expected to interact with the CSM gathered
through former mass-loss events. This would generate luminous
light curves, which may also show large luminosity fluctuations.
Even though this scenario could reproduce the observable proper-
ties of SN 2020wnt, we do not have a model to test it.
6 DISCUSSION
Similar pre-peak bump light curves have been detected in several
SNe, which include SLSNe-I SN 2006oz, LSQ14bdq, DES14X3taz,
DES15S2nr, DES17X1amf, SN 2018bsz and some signs of it
in SN 2018hti (Fiore et al. 2022), and the peculiar SNe Ib/c
SN 2005bf (Anupama et al. 2005; Tominaga et al. 2005; Folatelli
et al. 2006; Maeda et al. 2007), PTF11mnb (Taddia et al. 2018a),
and SN 2019cad (Gutiérrez et al. 2021). These bumps have been
suggested to be a common characteristic of the SLSN class (Nicholl
et al. 2016b), although more recent analyses with larger samples de-
termined that such feature is not ubiquitous to all SLSNe-I (Angus
et al. 2019).
Based on the analysis and comparison presented in previous
sections, we found that SN 2020wnt is an object that seems to
connect SN Ic and SLSN-I events. Although this connection was
12 https://github.com/manolis07gr/TigerFit
established before (e.g., Pastorello et al. 2010), the excellent pho-
tometric and spectroscopic coverage of this SN provides important
insights to better understand such relations, and the implications
for the explosion and progenitor stars. We discuss below the conse-
quences of these results.
6.1 Unusual light curve evolution
As discussed in previous sections, the light curves of SN 2020wnt
show remarkable features such as a pre-peak bump, a tail resembling
the 56Co decay followed by a sudden drop and minor luminosity
fluctuations. Bumpy light curves have been observed before max-
imum in several SLSNe-I, as shown in Figure 10, and have been
suggested to be frequent in the SLSN class (Nicholl et al. 2016b).
However, more recent analyses with larger samples determined that
such feature is not ubiquitous in SLSN-I (Angus et al. 2019). Early
bumps have been interpreted as resulting from the recombination
wave in the ejecta (Leloudas et al. 2012, for SN 2006oz), the shock
breakout within a dense CSM (Moriya & Maeda 2012), the shock
cooling of extended material around the progenitor (Nicholl et al.
2015b; Piro 2015; Smith et al. 2016; Angus et al. 2019) or an en-
hanced magnetar-driven shock breakout (Kasen et al. 2016). In the
case of SN 2020wnt, we propose that such a feature is a conse-
quence of a post-shock cooling phase in an extended progenitor
(Section 5.2).
On the other hand, tails resembling the 56Co decay have also
been previously observed in other SLSNe-I, however, it was sug-
gested that theymay be powered bymagnetar energy injection rather
than 56Co decay (Inserra et al. 2013, but see Gal-Yam 2012). In the
case of SN 2020wnt, and as shown in Section 5.2, the 56Ni models
provide a better representation of the light curve of SN 2020wnt than
those obtained with the magnetar model, suggesting that the main
source of power is radioactivity. Signs of CSM interaction (e.g.
presence of H lines in the spectra, light curve fluctuations) may
imply additional energy contribution from the CSM interaction.
Starting from 273 days from explosion, SN 2020wnt shows a
sudden drop in brightness. Although resembling breaks have been
observed in a few SLSNe I (e.g. Inserra et al. 2017), several inter-
acting objects (e.g. Mattila et al. 2008; Pastorello et al. 2008; Ofek
et al. 2014; Tartaglia et al. 2020), and in some Super-Chandrasekhar
SNe Ia (e.g. Hsiao et al. 2020), the slope measured in SN 2020wnt
is unique. While for slow-evolving SLSNe I, Inserra et al. (2017)
measured a decline that follows a power law of 𝑡−5, interacting
objects and Super-Chandrasekhar SNe Ia usually show less steep
declines (e.g. Fransson et al. 2014; Moriya 2014). On the contrary,
for SN 2020wnt, we find a much steeper slope, which follows a
power law of 𝑡−11.
It has been discussed that this break can either be caused by
the breakout of the shock through the dense shell (Ofek et al. 2014;
Fransson et al. 2014), or the result of a transition to a momentum-
conserving phase, occurring when shock runs over a mass of CSM
equivalent to the ejecta mass (Ofek et al. 2014, but see Moriya
2014), or the result of CO formation (Hsiao et al. 2020), or even
the consequence of dust formation in a cool, dense shell (Pastorello
et al. 2008; Mattila et al. 2008; Smith et al. 2008). Despite the
breakout of the shock through the dense shell being a plausible
explanation for a type IIn event such as SN 2010jl, it is less reliable
for SN 2020wnt. While SN 2020wnt shows some signatures of
interaction, its spectra lack flat-topped profiles and intermediate-
width features, which are expected from the interaction of the ejecta
with dense CSM. However, as mentioned in Section 5, ejecta-CSM
interaction may play an important role in the sudden drop observed
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 17
in the late optical light curve of SN 2020wnt. Although CO and dust
formation could be an option, it is not supported by the spectrum at
293 days from explosion (taken during the early part of the sudden
drop). Here, we do not detect any emission line blueshifted, typically
seen in SN spectra with dust forming in the ejecta. However, to
confirm or reject this alternative, NIR observations are needed.
6.2 Absence of O ii lines and presence of C ii lines
One of themost common features observed in the early phase spectra
of SLSNe-I is the presence of O ii lines around 4000 – 4500 Å
(Quimby et al. 2011, 2018; but see, Könyves-Tóth & Vinkó 2021).
These lines appear at high temperatures (12000-15000 K; Inserra
2019) as a consequence of non-thermal excitation (Mazzali et al.
2016). In the case of SN 2020wnt, we noticed that these lines are
not visible in the spectra at any epoch (see Figure 4). As shown in
Section 3.3 (and Figure 3), the temperature of SN 2020wnt evolves
in a range of values lower than 10000 K. This temperature is indeed
not high enough to ionise the oxygen.Quimby et al. (2018) suggested
that the lack of these features, at least in standard luminosity SNe Ic,
may be either the product of rapid cooling or due to a lack of non-
thermal sources of excitation. Given that the initial emission can be
reproduced as the results of the cooling of an extended progenitor,
this is not compatible with a rapid cooling phase. Therefore, the lack
of a non-thermal source seems to be the most reliable explanation.
This is also supported by our modelling, which does not require a
very extended mixing of radioactive material.
Gal-Yam (2019) found that, in addition to the O lines, the C
lines are also typical of SLSNe I at around peak, and suggested that
they result from the emission of an almost pure C/O envelope, with-
out significant contamination of higher-mass elements from deeper
layers. Inspecting the spectra of SN 2020wnt, we detect strong and
persistent lines of C ii. These lines are predicted by spectral mod-
els (Dessart et al. 2012; Mazzali et al. 2016; Dessart 2019), and
are identified in several SLSNe-I (e.g., PTF09cnd (Quimby et al.
2018), PTF10aagc (Quimby et al. 2018), PTF12dam, SN 2015bn,
Gaia16apd (Yan et al. 2017a), iPFT16bad (Yan et al. 2017b),
SN 2017gci, SN 2018bsz, SN 2018hti (Lin et al. 2020; Fiore et al.
2022), although inmost of the cases, the strength of the lines is mod-
erate. However, SN 2020wnt along with SN 2018bsz, have stronger
lines than these other objects, and furthermore their C ii lines agree
with the models of Dessart (2019), which typically overestimate the
observed strength of these lines.
In the spectra of SN 2020wnt, we also detect a strong Si ii. A
weak Si ii feature has been detected in a few objects (e.g. PTF09cnd
and SN 2015bn). From the modelling side, Si ii is not predicted by
the models presented in Mazzali et al. (2016), but it is reproduced
by the magnetar model presented in Dessart et al. (2012). They
argue that the presence of the Si ii, C ii, and He i lines is a result
of the extra energy from a magnetar that heats the material and
thermally excites the gas. Although the detection of these lines
in SN 2020wnt supports this scenario, our light curve modelling
(Section 5.2) disfavours it.
6.3 SN 2020wnt: an extreme case of SN 2007gr?
As discussed before (Section 4.3), the spectral evolution of
SN 2020wnt resembles that of the carbon-rich type Ic SN 2007gr.
In Section 4, we showed that their light curves evolve differently.
Indeed, SN 2020wnt is over 3 magnitudes brighter than SN 2007gr,
it evolves much more slowly (it is ∼ 4 times slower) and has the
pre-peak bump that was not detected in its fainter counterpart. How-
ever, the remarkable similarity of the spectra suggests that they may
have a progenitor with a similar composition. The question that
arises from this comparison is why SN 2020wnt and SN 2007gr
show comparable spectroscopic behaviour but completely different
photometric properties?
In our attempt to provide an answer, we analyse the en-
vironments of these objects. In Section 3.1 we mentioned that
SN 2020wnt is in a metal-poor environment (12 + log(O/H)= 8.175
dex by using the mass-metallicity relation of Kewley & Ellison
2008). More precisely, its host galaxy is faint, has a low stellar
mass and very little star formation, similar to those observed for
SLSNe-I. On the other hand, the host of SN 2007gr was identified
as a nearby spiral galaxy (M𝐵 = −18.8913; Makarov et al. 2014)
that also hosted SN 1969L and SN 1961V (Hunter et al. 2009;
Chen et al. 2014). According to Maund & Ramirez-Ruiz (2016),
SN 2007gr was located at near the centre of a dense young, massive
star association (Kuncarayakti et al. 2013). Modjaz et al. (2011)
measured the metallicity near the position of SN 2007gr by us-
ing the O3N2 diagnostic method (Pettini & Pagel 2004) and found
it was 12 + log(O/H)= 8.64 dex, indicating a metal-rich environ-
ment. These results lead us to conclude that the host environments
of SN 2020wnt and SN 2007gr are different, and based on these
estimations, also the metallicity of their progenitor stars.
The nature of the progenitor star and explosion of SN 2007gr
has been broadly discussed (e.g., Crockett et al. 2008; Chen et al.
2014; Maund & Ramirez-Ruiz 2016). Chen et al. (2014) suggested
that the progenitor of SN 2007gr was a low-mass Wolf–Rayet star
resulting from an interacting binary.Maund&Ramirez-Ruiz (2016)
support aWolf–Rayet progenitor, but add that it was an initiallymas-
sive star. From the nebular modelling, Mazzali et al. (2010) found
that SN 2007gr was the explosion of a low-mass CO core, probably
the result of a star with an initial mass of 15 M . For SN 2020wnt,
our light curve modelling suggests a massive progenitor and an en-
ergetic explosion with lot of 56Ni produced. These parameters are,
in all respects, more extreme than those found for SN 2007gr.
7 CONCLUSIONS
SN 2020wnt is a slow-evolving carbon-rich SLSN-I. Its light curves
show an early bump lasting ∼ 5 days followed by a slow rise to the
main peak. The peak is reached at different times, occurring faster
in the bluer bands. With an absolute peak magnitudes of around
∼ −20.5 mag, SN 2020wnt is in the low end of the luminosity
distribution of SLSNe-I. After 130 days from explosion, the light
curves show a linear decline in all bands, with slopes being around
the expected decline rate of the 56Co decay. Later, from 273 days,
a sudden drop in brightness is observed, implying a significant
leakage of gamma-ray photons. Our last observations (after 350
days from explosion), show an increase in brightness, which may
suggest interaction between the ejecta and the CSM. Indeed, minor
light curve fluctuations support this scenario.
During the photospheric phase, the optical spectra show clear
lines of C ii and Si ii, while the classical O ii lines that typically
characterise SLSNe-I are not detected. The lack of O ii lines is
probably related to the low temperatures of this object (below 10000
K). Late-time spectra display strong lines of [O i], [Ca ii], Ca ii,
Mg i], as well as, a broad emission of [O iii] and Balmer lines.
13 http://leda.univ-lyon1.fr/
MNRAS 000, 1–24 (2022)
18 Gutiérrez et al.
We modelled the light curve and the expansion velocities of
SN 2020wnt using a one-dimensional hydrodynamics code. Two
scenarios were investigated, with the radioactive nickel and the
magnetar as primary powering sources. In both cases, we found that
an extended progenitor was required to reproduce the time scale of
the peaks. However, the magnetar model produces a much worse
fit to the data. Therefore, we consider the 56Ni as the main power
source. Specifically, we found that SN 2020wnt can be explained
by a progenitor with a pre-SN mass of 80 M , a pre-SN radius
of 15 R , an explosion energy of 45 × 1051 erg, and ejecting 4
M of 56Ni. In this scenario, the first peak results from a post-
shock cooling phase for the extended progenitor, and the luminous
main peak is due to a large 56Ni production. The values of the
parameters obtained are consistent with those expected for a PISN,
which provide support for this scenario in the case of SN 2020wnt.
Although our model reproduces the almost complete evolution of
the light curve reasonably well, it fails to explain the excess of flux
at ∼ 245 days and the shape of the light curve after that. We propose
that this behaviour is probably due to an additional contribution of
ejecta-CSM interaction. The drop in brightness after 273 days from
explosion could be attributed to either the end of earlier ejecta-CSM
interaction, or the formation of molecules and dust in the SN ejecta
or in a shocked cool dense shell. However, NIR observation are
needed to confirm this suggestion.
We noticed remarkable spectral similarities between
SN 2020wnt and carbon-rich type Ic SN 2007gr, but on a longer
time scale. This resemblance may suggest a connection between
these two events, most probably associated with the structures of
their progenitor stars. Although we have found a model that can
explain the main photometric properties of SN 2020wnt, we also
discussed a possibility where the CSM interaction is the primary
energy source already at early phases. However, we did not explore
in detail this alternative scenario due to the lack of an available
model with such extreme input parameters as those observed in
SN 2020wnt.
ACKNOWLEDGEMENTS
We thank the anonymous referee for the comments and sugges-
tions that have helped to improve the paper. M.B. and M.O. ac-
knowledge support from UNRN PI2020 40B885 and grant PICT-
2020-SERIEA-01141 and PIP 112-202001-10034. A.R. acknowl-
edges support from ANID BECAS/DOCTORADO NACIONAL
21202412. T.M.R. acknowledges the financial support of the Finnish
Academy of Science and Letters. N.E.R. acknowledges partial sup-
port from MIUR, PRIN 2017 (grant 20179ZF5KS), from the Span-
ish MICINN grant PID2019-108709GB-I00 and FEDER funds,
and from the program Unidad de Excelencia María de Maeztu
CEX2020-001058-M. M.F. is supported by a Royal Society - Sci-
ence Foundation Ireland University Research Fellowship. S.M. ac-
knowledges support from theMagnus Ehrnrooth Foundation and the
Vilho, Yrjö and Kalle Väisälä Foundation. M.S. was a visiting as-
tronomer at the Infrared Telescope Facility, which is operated by the
University of Hawaii under contract 80HQTR19D0030with the Na-
tional Aeronautics and Space Administration. X.W. is supported by
the National Science Foundation of China (NSFC grants 12033003
and 11633002), the Scholar Program of Beijing Academy of Sci-
ence and Technology (DZ:BS202002), and the Tencent Xplorer
Prize. J.Z. is supported by the NSFC (grants 12173082, 11773067),
by the Youth Innovation Promotion Association of the CAS (grant
2018081), and by the Ten Thousand Talents Program of Yunnan for
Top-notch Young Talents. Funding for the LJT has been provided
by Chinese Academy of Sciences and the People’s Government of
Yunnan Province. The LJT is jointly operated and administrated by
Yunnan Observatories and Center for Astronomical Mega-Science,
CAS. This work is supported by the National Natural Science Foun-
dation of China (NSFC grants 12033003, 11633002, 11325313,
and 11761141001), the National Program on Key Research and
Development Project (grant no. 2016YFA0400803). This work is
partially supported by China Manned Spaced Project (CMS-CSST-
2021-A12). This work is funded by China Postdoctoral Science
Foundation (grant no. 2021M691821).
We are grateful to Jan Aaltonen, Ville Antila, Katja Matilainen and
Sofia Suutarinen, who observed this target as part of the "NOT
course 2021" organised by the Department of Physics and Astron-
omy at the University of Turku in October 2021.
Based on observations made with the Nordic Optical Telescope,
owned in collaboration by the University of Turku and Aarhus Uni-
versity, and operated jointly by Aarhus University, the University of
Turku and the University of Oslo, representing Denmark, Finland
and Norway, the University of Iceland and Stockholm University at
the Observatorio del Roque de los Muchachos, La Palma, Spain, of
the Instituto de Astrofisica de Canarias.
Observations from the NOT were obtained through the NUTS2 col-
laboration which are supported in part by the Instrument Centre
for Danish Astrophysics (IDA). The data presented here were ob-
tained in part with ALFOSC, which is provided by the Instituto de
Astrofisica de Andalucia (IAA) under a joint agreement with the
University of Copenhagen and NOTSA.
Partially based on observations collected at Copernico and Schmidt
telescopes (Asiago, Italy) of the INAF - Osservatorio Astronomico
di Padova.
Based on observations made with the Gran Telescopio Canarias
(GTC), installed in the Spanish Observatorio del Roque de los
Muchachos of the Instituto de Astrofísica de Canarias, in the is-
land of La Palma.
We acknowledge the support of the staff of the Lijiang 2.4m, Xing-
long 2.16m and 80cm telescopes.
The Liverpool Telescope is operated on the island of La Palma by
Liverpool John Moores University in the Spanish Observatorio del
Roque de los Muchachos of the Instituto de Astrofisica de Canarias
with financial support from the UK Science and Technology Facil-
ities Council.
This work hasmade use of data from the Asteroid Terrestrial-impact
Last Alert System (ATLAS) project. The Asteroid Terrestrial-
impact Last Alert System (ATLAS) project is primarily funded to
search for near earth asteroids through NASA grants NN12AR55G,
80NSSC18K0284, and 80NSSC18K1575; byproducts of the NEO
search include images and catalogs from the survey area. This work
was partially funded by Kepler/K2 grant J1944/80NSSC19K0112
and HST GO-15889, and STFC grants ST/T000198/1 and
ST/S006109/1. The ATLAS science products have been made pos-
sible through the contributions of the University of Hawaii Institute
for Astronomy, the Queen’s University Belfast, the Space Telescope
Science Institute, the South African Astronomical Observatory, and
The Millennium Institute of Astrophysics (MAS), Chile.
The ZTF forced-photometry service was funded under the Heising-
Simons Foundation grant #12540303 (PI: Graham).
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 19
DATA AVAILABILITY
The data underlying this article are available in the appendix A
(Tables A1 – A5) and through the WISeREP (https://wiserep.
weizmann.ac.il/home) archive (Yaron & Gal-Yam 2012).
REFERENCES
Ambikasaran S., et al., 2016, IEEE Trans. Pattern Anal. Mach. Intell., 38
Anderson J. P., et al., 2018, A&A, 620, A67
Angus C. R., et al., 2019, MNRAS, 487, 2215
Anupama G. C., Sahu D. K., Deng J., Nomoto K., Tominaga N., Tanaka M.,
Mazzali P. A., Prabhu T. P., 2005, ApJ, 631, L125
Astropy Collaboration et al., 2018, AJ, 156, 123
Barkat Z., Rakavy G., Sack N., 1967, Phys. Rev. Lett., 18, 379
Bellm E. C., et al., 2019, PASP, 131, 018002
Bersten M. C., Benvenuto O., Hamuy M., 2011, ApJ, 729, 61
Bersten M. C., Benvenuto O. G., Orellana M., Nomoto K., 2016, ApJ, 817,
L8
Bianco F. B., et al., 2014, ApJS, 213, 19
Bradley L., et al., 2019, astropy/photutils: v0.6,
doi:10.5281/zenodo.2533376, https://doi.org/10.5281/
zenodo.2533376
Cappellaro E., 2014, SNOoPY: a package for SN photometry, http://
sngroup.oapd.inaf.it/snoopy.html
Cardelli J. A., Clayton G. C., Mathis J. S., 1989, ApJ, 345, 245
Chambers K. C., et al., 2016, preprint, (arXiv:1612.05560)
Chatzopoulos E., et al., 2011, ArXiv e-prints:1101.3581,
Chatzopoulos E., Wheeler J. C., Vinko J., Nagy A. P., Wiggins B. K., Even
W. P., 2016, ApJ, 828, 94
Chen T.-W., et al., 2013, ApJ, 763, L28
Chen J., et al., 2014, ApJ, 790, 120
Chen T. W., et al., 2015, MNRAS, 452, 1567
Chen T.-W., Smartt S. J., Yates R. M., Nicholl M., Krühler T., Schady P.,
Dennefeld M., Inserra C., 2017, MNRAS, 470, 3566
Chen T. W., et al., 2021, arXiv e-prints, p. arXiv:2109.07942
Chevalier R. A., Irwin C. M., 2011, ApJ, 729, L6
Childress M., et al., 2013, ApJ, 770, 107
Chonis T. S., Gaskell C. M., 2008, AJ, 135, 264
Crockett R. M., et al., 2008, ApJ, 672, L99
Cushing M. C., Vacca W. D., Rayner J. T., 2004, PASP, 116, 362
De Cia A., et al., 2018, ApJ, 860, 100
Dessart L., 2019, A&A, 621, A141
Dessart L., Hillier D. J., Waldman R., Livne E., Blondin S., 2012, MNRAS,
426, L76
Dessart L., Waldman R., Livne E., Hillier D. J., Blondin S., 2013, MNRAS,
428, 3227
Dessart L., Audit E., Hillier D. J., 2015, MNRAS, 449, 4304
Dessart L., Hillier D. J., Sukhbold T., Woosley S. E., Janka H. T., 2021,
A&A, 652, A64
Elmhamdi A., Danziger I. J., Cappellaro E., Della Valle M., Gouiffes C.,
Phillips M. M., Turatto M., 2004, A&A, 426, 963
Fesen R. A., et al., 1999, AJ, 117, 725
Fioc M., Rocca-Volmerange B., 1997, A&A, 326, 950
Fiore A., et al., 2021, MNRAS, 502, 2120
Fiore A., et al., 2022, MNRAS, 512, 4484
Fisher A. K., 2000, PhD thesis, THE UNIVERSITY OF OKLAHOMA
Folatelli G., et al., 2006, ApJ, 641, 1039
Förster F., et al., 2021, AJ, 161, 242
Fransson C., Chevalier R. A., 1989, ApJ, 343, 323
Fransson C., et al., 2002, ApJ, 572, 350
Fransson C., et al., 2014, ApJ, 797, 118
Gal-Yam A., 2012, Science, 337, 927
Gal-Yam A., 2019, ApJ, 882, 102
Gal-Yam A., et al., 2009, Nature, 462, 624
Ginzburg S., Balberg S., 2012, ApJ, 757, 178
Gomez S., Berger E., Nicholl M., Blanchard P. K., Hosseinzadeh G., 2022,
arXiv e-prints, p. arXiv:2204.08486
Graham M. J., et al., 2019, PASP, 131, 078001
Gutiérrez C. P., et al., 2017, ApJ, 850, 89
Gutiérrez C. P., et al., 2020a, MNRAS, 496, 95
Gutiérrez C. P., et al., 2020b, MNRAS, 499, 974
Gutiérrez C. P., et al., 2021, MNRAS, 504, 4907
Harutyunyan A. H., et al., 2008, A&A, 488, 383
Heger A., Woosley S. E., 2002, ApJ, 567, 532
Heger A., Fryer C. L., Woosley S. E., Langer N., Hartmann D. H., 2003,
ApJ, 591, 288
Houck J. C., Fransson C., 1996, ApJ, 456, 811
Hsiao E. Y., et al., 2019, PASP, 131, 014002
Hsiao E. Y., et al., 2020, ApJ, 900, 140
Huang F., Li J.-Z., Wang X.-F., Shang R.-C., Zhang T.-M., Hu J.-Y., Qiu
Y.-L., Jiang X.-J., 2012, Research in Astronomy and Astrophysics, 12,
1585
Hunter D. J., et al., 2009, A&A, 508, 371
Inserra C., 2019, Nature Astronomy, 3, 697
Inserra C., et al., 2013, ApJ, 770, 128
Inserra C., et al., 2017, MNRAS, 468, 4642
Inserra C., et al., 2018, MNRAS, 475, 1046
Janka H.-T., 2012, Annual Review of Nuclear and Particle Science, 62, 407
Jerkstrand A., 2017, Spectra of Supernovae in the Nebular Phase. p. 795,
doi:10.1007/978-3-319-21846-5_29
Jerkstrand A., Smartt S. J., Fraser M., Fransson C., Sollerman J., Taddia F.,
Kotak R., 2014, MNRAS, 439, 3694
Jerkstrand A., Smartt S. J., Heger A., 2016, MNRAS, 455, 3207
Jerkstrand A., et al., 2017, ApJ, 835, 13
Kangas T., et al., 2017, MNRAS, 469, 1246
Kaplan N., Soker N., 2020, MNRAS, 492, 3013
Kasen D., 2010, ApJ, 708, 1025
Kasen D., Bildsten L., 2010, ApJ, 717, 245
Kasen D., Metzger B. D., Bildsten L., 2016, ApJ, 821, 36
Kewley L. J., Ellison S. L., 2008, ApJ, 681, 1183
Könyves-Tóth R., Vinkó J., 2021, ApJ, 909, 24
Kozyreva A., Yoon S. C., Langer N., 2014, A&A, 566, A146
Kumar A., et al., 2020, ApJ, 892, 28
Kuncarayakti H., et al., 2013, AJ, 146, 30
Kuncarayakti H., et al., 2015, A&A, 579, A95
Le Borgne D., Rocca-Volmerange B., 2002, A&A, 386, 446
Leloudas G., et al., 2012, A&A, 541, A129
Leloudas G., et al., 2015, MNRAS, 449, 917
Li H., McCray R., 1993, ApJ, 405, 730
Lin W. L., et al., 2020, MNRAS, 497, 318
Lunnan R., et al., 2014, ApJ, 787, 138
Lunnan R., et al., 2016, ApJ, 831, 144
Maeda K., Mazzali P. A., Deng J., Nomoto K., Yoshii Y., Tomita H.,
Kobayashi Y., 2003, ApJ, 593, 931
Maeda K., et al., 2007, ApJ, 666, 1069
Magnier E. A., et al., 2020, ApJS, 251, 6
Makarov D., Prugniel P., Terekhova N., Courtois H., Vauglin I., 2014, A&A,
570, A13
Masci F. J., et al., 2019, PASP, 131, 018003
Mattila S., et al., 2008, MNRAS, 389, 141
Maund J. R., Ramirez-Ruiz E., 2016, MNRAS, 456, 3175
Mazzali P. A., Maurer I., Valenti S., Kotak R., Hunter D., 2010, MNRAS,
408, 87
Mazzali P. A., Sullivan M., Pian E., Greiner J., Kann D. A., 2016, MNRAS,
458, 3455
Mazzali P. A., Moriya T. J., Tanaka M., Woosley S. E., 2019, MNRAS, 484,
3451
Milisavljevic D., Fesen R. A., Chevalier R. A., Kirshner R. P., Challis P.,
Turatto M., 2012, ApJ, 751, 25
Millard J., et al., 1999, ApJ, 527, 746
Modjaz M., Kewley L., Bloom J. S., Filippenko A. V., Perley D., Silverman
J. M., 2011, ApJ, 731, L4+
Moriya T. J., 2014, arXiv e-prints, p. arXiv:1402.2519
MNRAS 000, 1–24 (2022)
20 Gutiérrez et al.
Moriya T. J., Maeda K., 2012, ApJ, 756, L22
Moriya T. J., Sorokina E. I., Chevalier R. A., 2018, Space Sci. Rev., 214, 59
Moriya T. J.,MuraseK., KashiyamaK., Blinnikov S. I., 2022, arXiv e-prints,
p. arXiv:2202.03082
Nasa High Energy Astrophysics Science Archive Research Center
(Heasarc) 2014, HEAsoft: Unified Release of FTOOLS and XANADU
(ascl:1408.004)
Neill J. D., et al., 2011, ApJ, 727, 15
Nicholl M., 2018, Research Notes of the American Astronomical Society,
2, 230
Nicholl M., 2021, Astronomy and Geophysics, 62, 5.34
Nicholl M., et al., 2013, Nature, 502, 346
Nicholl M., et al., 2015a, MNRAS, 452, 3869
Nicholl M., et al., 2015b, ApJ, 807, L18
Nicholl M., et al., 2016a, ApJ, 826, 39
Nicholl M., et al., 2016b, ApJ, 828, L18
Ofek E. O., et al., 2014, ApJ, 781, 42
Orellana M., Bersten M. C., Moriya T. J., 2018, A&A, 619, A145
Pastorello A., et al., 2008, MNRAS, 389, 113
Pastorello A., et al., 2010, ApJ, 724, L16
Perley D. A., et al., 2016, ApJ, 830, 13
Pettini M., Pagel B. E. J., 2004, MNRAS, 348, L59
Piro A. L., 2015, ApJ, 808, L51
Prentice S. J., et al., 2019, MNRAS, 485, 1559
Pursiainen M., et al., 2022, arXiv e-prints, p. arXiv:2202.01635
Quimby R. M., et al., 2011, Nature, 474, 487
Quimby R. M., et al., 2018, ApJ, 855, 2
Rakavy G., Shaviv G., 1967, ApJ, 148, 803
Schlafly E. F., Finkbeiner D. P., 2011, ApJ, 737, 103
Schulze S., et al., 2018, MNRAS, 473, 1258
Skrutskie M. F., et al., 2006, AJ, 131, 1163
Smith N., Foley R. J., Filippenko A. V., 2008, ApJ, 680, 568
Smith M., et al., 2016, ApJ, 818, L8
Smith K. W., et al., 2020, arXiv e-prints, p. arXiv:2003.09052
Soker N., 2022, arXiv e-prints, p. arXiv:2205.09560
Soker N., Gilkis A., 2017, ApJ, 851, 95
Sollerman J., Kozma C., Fransson C., Leibundgut B., Lundqvist P., Ryde F.,
Woudt P., 2000, ApJ, 537, L127
Sorokina E., Blinnikov S., Nomoto K., Quimby R., Tolstov A., 2016, ApJ,
829, 17
Stoll R., Prieto J. L., Stanek K. Z., Pogge R. W., Szczygiel D. M., Pojmanski
G., Antognini J., Yan H., 2011, ApJ, 730, 34
Stritzinger M. D., et al., 2018, A&A, 609, A135
Sullivan M., et al., 2010, MNRAS, 406, 782
Taddia F., et al., 2018a, A&A, 609, A106
Taddia F., et al., 2018b, A&A, 609, A136
Tartaglia L., et al., 2020, A&A, 635, A39
Taubenberger S., et al., 2006, MNRAS, 371, 1459
Terreran G., et al., 2017, Nature Astronomy, 1, 713
Terreran G., et al., 2019, ApJ, 883, 147
Tinyanont S., Dimitriadis G., Foley R. J., 2020, Transient Name Server
Classification Report, 2020-3473, 1
Tominaga N., et al., 2005, ApJ, 633, L97
Tonry J. L., et al., 2018, Publications of the Astronomical Society of the
Pacific, 130, 064505
Tremonti C. A., et al., 2004, ApJ, 613, 898
Tully R. B., et al., 2013, AJ, 146, 86
Valenti S., et al., 2008a, MNRAS, 383, 1485
Valenti S., et al., 2008b, ApJ, 673, L155
Vreeswijk P. M., et al., 2017, ApJ, 835, 58
Vurm I., Metzger B. D., 2021, ApJ, 917, 77
Wheeler J. C., Chatzopoulos E., Vinkó J., Tuminello R., 2017, ApJ, 851,
L14
Wilk K. D., Hillier D. J., Dessart L., 2019, MNRAS, 487, 1218
Woosley S. E., 2010, ApJ, 719, L204
Woosley S. E., Hartmann D., Pinto P. A., 1989, ApJ, 346, 395
Yan L., et al., 2017a, ApJ, 840, 57
Yan L., et al., 2017b, ApJ, 848, 6
Yaron O., Gal-Yam A., 2012, PASP, 124, 668
Young D. R., et al., 2010, A&A, 512, A70+
SUPPORTING INFORMATION
Supplementary data are available at MNRAS online.
Table A1. Optical photometry of SN 2020wnt.
Table A2. JHK Vega photometry of SN 2020wnt obtained with
NOTCam.
Table A3. UV photometry obtained with Swift in the AB system.
Table A4. ATLAS AB optical photometry.
Table A5. ZTF AB optical photometry.
Table A6. Spectroscopic observations of SN 2020wnt.
APPENDIX A: TABLES
AFFILIATIONS
1 Finnish Centre for Astronomy with ESO (FINCA), FI-20014
University of Turku, Finland
2 Tuorla Observatory, Department of Physics and Astronomy,
FI-20014 University of Turku, Finland
3 INAF - Osservatorio Astronomico di Padova, Vicolo
dell’Osservatorio 5, I-35122 Padova, Italy
4 Facultad de Ciencias Astronómicas y Geofísicas, Universidad
Nacional de La Plata, Paseo del Bosque S/N, B1900FWA,
La Plata, Argentina
5 Instituto de Astrofísica de La Plata (IALP), CCT-CONICET-
UNLP. Paseo del Bosque S/N, B1900FWA, La Plata, Argentina
6 Kavli Institute for the Physics and Mathematics of the Universe
(WPI), The University of Tokyo, 5-1-5 Kashiwanoha,
Kashiwa, Chiba 277-8583, Japan
7 Universidad Nacional de Río Negro. Sede Andina, Mitre 630
(8400), Bariloche, Argentina.
8 Consejo Nacional de Investigaciones Científicas y Tećnicas
(CONICET), Argentina.
9 European Centre for Theoretical Studies in Nuclear Physics and
Related Areas (ECT∗), Fondazione Bruno Kessler, Trento, Italy
10 INFN-TIFPA, Trento Institute for Fundamental Physics and
Applications, Via Sommarive 14, I-38123 Trento, Italy
11 Department of Physics and Astronomy, Aarhus University, Ny
Munkegade 120, DK-8000 Aarhus C, Denmark
12 Departamento de Ciencias Físicas – Universidad Andrés Bello,
Avda. República 252, Santiago, Chile
13 Millennium Institute of Astrophysics, Nuncio Monsenor Sótero
Sanz 100, Providencia, Santiago, Chile
14 Cosmic Dawn Center (DAWN), Denmark
15 Niels Bohr Institute, University of Copenhagen, Jagtvej 128,
2200 Copenhagen N, Denmark
16 Department of Physics and Astronomy G. Galilei, University of
Padova, Vicolo dell’Osservatorio 3, 35122, Padova, Italy
17 Astrophysics Research Institute, Liverpool John Moores Univer-
sity, IC2, Liverpool Science Park, 146 Brownlow Hill, Liverpool
L3 5RF, UK
18 Max-Planck-Institut für Astrophysik, Karl-Schwarzschild Str. 1,
D-85748 Garching, Germany
19 School of Physics and Astronomy, University of Southampton,
Southampton, SO17 1BJ, UK
20 Physics Department and Tsinghua Center for Astrophysics
(THCA), Tsinghua University, Beijing 100084, China
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 21
Table A1. Optical photometry of SN 2020wnt.
UT date MJD Phase 𝑢 𝐵 𝑉 𝑔 𝑟 𝑖 𝑧 Telescope
(days)∗ (mag) (mag) (mag) (mag) (mag) (mag) (mag)
20201121 59174.83 38.21 18.43 ± 0.14 17.91 ± 0.05 17.28 ± 0.03 17.57 ± 0.03 17.11 ± 0.03 16.94 ± 0.06 – Schmidt Telescope
20201121 59174.99 38.36 18.40 ± 0.03 17.93 ± 0.02 17.25 ± 0.01 17.50 ± 0.01 17.07 ± 0.01 16.92 ± 0.01 16.87 ± 0.02 Copernico 1.82-m
20201124 59177.13 40.44 18.38 ± 0.10 17.84 ± 0.04 17.16 ± 0.03 17.44 ± 0.03 17.01 ± 0.03 16.84 ± 0.04 – Schmidt Telescope
20201126 59179.84 43.06 – 17.67 ± 0.06 17.05 ± 0.06 17.37 ± 0.07 16.86 ± 0.07 16.73 ± 0.07 – Schmidt Telescope
20201129 59182.83 45.96 – 17.63 ± 0.07 16.94 ± 0.06 17.28 ± 0.11 16.81 ± 0.05 16.69 ± 0.06 – Schmidt Telescope
20201203 59186.97 49.97 18.01 ± 0.02 17.44 ± 0.02 16.83 ± 0.01 17.08 ± 0.01 16.66 ± 0.01 16.49 ± 0.01 16.47 ± 0.01 LT
20201205 59188.01 50.98 17.98 ± 0.02 – – 17.03 ± 0.01 16.64 ± 0.01 16.46 ± 0.01 16.50 ± 0.02 LT
20201207 59190.97 53.85 17.89 ± 0.02 17.34 ± 0.01 16.72 ± 0.01 16.93 ± 0.01 16.54 ± 0.01 16.37 ± 0.01 16.35 ± 0.01 LT
20201210 59193.97 56.75 17.84 ± 0.02 17.28 ± 0.01 16.66 ± 0.01 16.87 ± 0.01 16.49 ± 0.01 16.28 ± 0.01 16.28 ± 0.01 LT
20201212 59195.97 58.69 17.80 ± 0.17 17.23 ± 0.03 16.68 ± 0.16 16.94 ± 0.05 16.46 ± 0.03 16.33 ± 0.03 16.22 ± 0.04 LT
20201213 59196.09 58.81 – 17.25 ± 0.14 16.67 ± 0.03 16.91 ± 0.03 16.47 ± 0.07 16.39 ± 0.07 – Schmidt Telescope
20201213 59196.97 59.66 17.82 ± 0.02 17.24 ± 0.01 16.60 ± 0.01 16.87 ± 0.01 16.44 ± 0.01 16.28 ± 0.03 16.24 ± 0.01 LT
20201214 59197.98 60.64 – 17.26 ± 0.04 16.66 ± 0.03 16.90 ± 0.03 16.47 ± 0.04 16.35 ± 0.04 – Schmidt Telescope
20201216 59199.88 62.48 – – – 16.86 ± 0.01 16.47 ± 0.01 16.33 ± 0.02 16.23 ± 0.02 Copernico 1.82-m
20201216 59199.95 62.55 – 17.24 ± 0.01 16.60 ± 0.01 – – – – LT
20201217 59200.95 63.52 17.82 ± 0.02 17.25 ± 0.01 16.58 ± 0.02 16.84 ± 0.01 16.41 ± 0.01 16.20 ± 0.01 16.20 ± 0.01 LT
20201218 59201.92 64.46 – 17.27 ± 0.05 16.65 ± 0.03 16.91 ± 0.02 16.46 ± 0.03 16.23 ± 0.05 – Schmidt Telescope
20201218 59201.94 64.48 17.83 ± 0.02 17.25 ± 0.02 16.60 ± 0.03 16.84 ± 0.01 16.43 ± 0.00 16.20 ± 0.01 16.18 ± 0.01 LT
20201220 59203.75 66.23 17.87 ± 0.03 17.24 ± 0.03 16.62 ± 0.02 16.88 ± 0.02 16.48 ± 0.03 16.24 ± 0.01 16.22 ± 0.02 Copernico 1.82-m
20201227 59210.04 72.33 18.04 ± 0.03 – – 16.87 ± 0.01 16.39 ± 0.01 16.18 ± 0.01 16.13 ± 0.01 NOT
20201229 59212.68 74.88 – – 16.69 ± 0.22 – 16.35 ± 0.06 16.03 ± 0.18 – TNT
20201230 59213.66 75.83 – – 16.62 ± 0.14 16.85 ± 0.06 16.35 ± 0.07 16.15 ± 0.12 – TNT
20201231 59214.69 76.83 – 17.31 ± 0.03 16.59 ± 0.10 16.83 ± 0.04 16.35 ± 0.05 16.18 ± 0.10 – TNT
20210101 59215.02 77.15 – 17.27 ± 0.15 16.62 ± 0.07 16.88 ± 0.06 16.43 ± 0.04 16.18 ± 0.03 16.19 ± 0.05 LT
20210101 59215.95 78.05 – – 16.64 ± 0.23 – 16.44 ± 0.16 16.18 ± 0.19 – LT
20210105 59219.47 81.46 – 17.42 ± 0.02 16.59 ± 0.01 16.94 ± 0.01 16.39 ± 0.05 16.20 ± 0.01 – TNT
20210109 59223.45 85.32 – 17.49 ± 0.03 16.60 ± 0.02 16.98 ± 0.04 16.40 ± 0.01 16.18 ± 0.01 – TNT
20210110 59224.96 86.78 – 17.51 ± 0.02 16.70 ± 0.02 17.09 ± 0.02 16.51 ± 0.02 16.18 ± 0.01 16.14 ± 0.01 Copernico 1.82-m
20210112 59226.65 88.42 – 17.66 ± 0.02 16.67 ± 0.03 17.12 ± 0.03 16.47 ± 0.02 16.25 ± 0.02 – TNT
20210113 59227.55 89.29 – 17.66 ± 0.02 16.70 ± 0.02 17.19 ± 0.02 16.43 ± 0.02 16.23 ± 0.02 – TNT
20210114 59228.05 89.78 – 17.67 ± 0.02 16.72 ± 0.02 17.15 ± 0.02 16.47 ± 0.02 16.20 ± 0.01 – Schmidt Telescope
20210115 59229.59 91.27 – 17.74 ± 0.02 16.75 ± 0.02 17.22 ± 0.03 16.50 ± 0.02 16.24 ± 0.02 – TNT
20210116 59230.63 92.28 – 17.84 ± 0.02 16.78 ± 0.02 17.25 ± 0.03 16.51 ± 0.02 16.26 ± 0.01 – TNT
20210117 59231.59 93.21 – 17.89 ± 0.02 16.82 ± 0.02 17.31 ± 0.02 16.54 ± 0.02 16.30 ± 0.02 – TNT
20210117 59231.86 93.47 – 17.83 ± 0.02 16.83 ± 0.02 17.30 ± 0.02 16.54 ± 0.03 16.23 ± 0.02 – Schmidt Telescope
20210118 59232.81 94.39 18.87 ± 0.04 17.84 ± 0.02 16.85 ± 0.02 17.35 ± 0.01 16.56 ± 0.01 16.28 ± 0.01 16.18 ± 0.01 Copernico 1.82-m
20210120 59234.86 96.38 19.06 ± 0.05 17.98 ± 0.02 16.85 ± 0.03 17.41 ± 0.02 16.52 ± 0.02 16.26 ± 0.04 16.11 ± 0.02 NOT
20210125 59239.95 101.31 19.41 ± 0.07 18.21 ± 0.02 16.95 ± 0.01 17.57 ± 0.01 16.57 ± 0.01 16.25 ± 0.01 16.15 ± 0.01 LT
20210128 59242.75 104.02 – 18.31 ± 0.06 17.01 ± 0.02 17.67 ± 0.03 16.66 ± 0.02 16.29 ± 0.02 – Schmidt Telescope
20210130 59244.57 105.78 – 18.40 ± 0.02 17.15 ± 0.02 17.78 ± 0.03 16.68 ± 0.02 16.40 ± 0.02 – TNT
20210130 59244.94 106.14 19.77 ± 0.10 18.40 ± 0.03 17.10 ± 0.01 17.76 ± 0.02 16.68 ± 0.01 16.31 ± 0.01 16.23 ± 0.01 LT
20210131 59245.56 106.74 – 18.33 ± 0.02 17.18 ± 0.02 17.81 ± 0.03 16.75 ± 0.02 16.44 ± 0.02 – TNT
20210201 59246.54 107.69 – 18.49 ± 0.02 17.16 ± 0.02 17.86 ± 0.03 16.76 ± 0.02 16.46 ± 0.02 – TNT
20210202 59247.58 108.70 – 18.56 ± 0.02 17.21 ± 0.02 17.92 ± 0.03 16.82 ± 0.02 16.50 ± 0.02 – TNT
20210204 59249.78 110.83 – 18.67 ± 0.01 17.39 ± 0.08 18.13 ± 0.13 16.90 ± 0.05 16.56 ± 0.03 – Schmidt Telescope
20210208 59253.86 114.79 – – 17.52 ± 0.02 18.31 ± 0.03 17.13 ± 0.03 16.76 ± 0.16 16.55 ± 0.03 Copernico 1.82-m
20210216 59261.87 122.55 21.19 ± 0.09 19.52 ± 0.02 17.89 ± 0.01 18.78 ± 0.01 17.44 ± 0.01 17.05 ± 0.01 16.76 ± 0.01 NOT
20210219 59264.48 125.08 – 19.49 ± 0.02 18.07 ± 0.02 18.75 ± 0.03 17.51 ± 0.02 17.09 ± 0.09 – TNT
20210219 59264.75 125.34 – – 18.15 ± 0.06 – 17.63 ± 0.06 17.14 ± 0.03 – Schmidt Telescope
20210220 59265.81 126.37 – 19.75 ± 0.15 – 18.89 ± 0.11 – – – Schmidt Telescope
20210221 59266.48 127.02 – 19.50 ± 0.02 18.17 ± 0.02 18.89 ± 0.03 17.55 ± 0.02 17.18 ± 0.02 – TNT
20210222 59267.46 127.97 – 19.67 ± 0.02 18.20 ± 0.02 19.08 ± 0.03 17.58 ± 0.02 17.25 ± 0.02 – TNT
20210226 59271.77 132.14 – – 18.30 ± 0.05 – 17.77 ± 0.04 17.33 ± 0.07 – Schmidt Telescope
20210228 59273.77 134.08 – – – 19.17 ± 0.13 – – – Schmidt Telescope
20210304 59277.53 137.72 – 19.74 ± 0.02 18.33 ± 0.07 19.13 ± 0.22 17.70 ± 0.02 17.26 ± 0.02 – TNT
20210306 59279.80 139.92 – – 18.37 ± 0.05 – 17.80 ± 0.04 17.39 ± 0.03 – Schmidt Telescope
20210313 59286.79 146.70 – – – – 17.90 ± 0.05 17.41 ± 0.02 17.06 ± 0.03 Copernico 1.82-m
20210318 59291.86 151.61 21.65 ± 0.25 20.00 ± 0.02 18.36 ± 0.01 19.19 ± 0.04 17.87 ± 0.01 17.41 ± 0.01 17.01 ± 0.01 NOT
20210323 59296.79 156.39 – – – – 17.96 ± 0.05 17.55 ± 0.04 – Schmidt Telescope
20210407 59311.49 170.63 – – – – 18.06 ± 0.25 17.74 ± 0.07 – TNT
20210416 59320.81 179.66 – – – – 18.18 ± 0.08 17.82 ± 0.05 – Schmidt Telescope
20210628 59393.07 249.68 – – – – – 18.46 ± 0.10 – Schmidt Telescope
20210630 59395.08 251.63 – – – 19.47 ± 0.27 – – – Schmidt Telescope
20210701 59396.06 252.58 – – – 19.44 ± 0.14 – – – Schmidt Telescope
20210706 59401.06 257.83 – 20.29 ± 0.19 – – – – – Schmidt Telescope
20210707 59402.07 257.42 – – 19.13 ± 0.10 – – – – Schmidt Telescope
20210708 59403.07 259.37 – – – 19.81 ± 0.18 – 18.55 ± 0.08 – Schmidt Telescope
20210717 59412.22 268.24 > 20.87 20.37 ± 0.04 19.05 ± 0.05 19.50 ± 0.08 18.91 ± 0.03 18.53 ± 0.07 17.84 ± 0.05 LT
20210722 59417.22 273.08 > 21.29 20.38 ± 0.04 19.22 ± 0.05 19.66 ± 0.06 18.97 ± 0.02 18.57 ± 0.11 17.97 ± 0.17 LT
20210729 59424.05 279.70 – – – – 19.38 ± 0.11 19.23 ± 0.10 – Schmidt Telescope
20210729 59424.18 279.83 > 20.97 20.88 ± 0.08 19.53 ± 0.08 19.84 ± 0.08 19.36 ± 0.09 19.05 ± 0.09 18.32 ± 0.07 LT
20210801 59427.21 282.76 > 21.66 21.15 ± 0.16 19.62 ± 0.04 20.16 ± 0.06 19.51 ± 0.06 19.18 ± 0.12 18.54 ± 0.13 LT
20210804 59430.18 285.64 – 21.00 ± 0.24 19.75 ± 0.06 20.31 ± 0.06 19.57 ± 0.09 19.22 ± 0.19 18.58 ± 0.13 LT
20210808 59434.15 289.49 – 21.18 ± 0.10 19.95 ± 0.10 20.48 ± 0.09 19.71 ± 0.11 19.43 ± 0.25 18.80 ± 0.17 LT
20210811 59437.16 292.40 – 21.19 ± 0.34 20.10 ± 0.11 20.63 ± 0.11 19.82 ± 0.07 19.62 ± 0.14 18.96 ± 0.18 LT
20210814 59440.17 295.32 – 21.56 ± 0.32 20.32 ± 0.38 20.68 ± 0.42 20.05 ± 0.59 19.71 ± 0.81 18.98 ± 0.41 LT
20210818 59444.18 299.21 – – 20.15 ± 0.13 20.67 ± 0.13 19.97 ± 0.09 19.73 ± 0.11 19.33 ± 0.10 LT
20210821 59447.19 302.12 – – 20.61 ± 0.15 21.07 ± 0.15 20.28 ± 0.15 19.87 ± 0.34 19.50 ± 0.21 LT
20210828 59454.10 308.82 – – 20.67 ± 0.14 21.10 ± 0.24 20.57 ± 0.11 20.08 ± 0.25 19.77 ± 0.17 LT
MNRAS 000, 1–24 (2022)
22 Gutiérrez et al.
Table A1. – continued
UT date MJD Phase 𝑢 𝐵 𝑉 𝑔 𝑟 𝑖 𝑧 Telescope
(days)∗ (mag) (mag) (mag) (mag) (mag) (mag) (mag)
20210831 59457.19 311.81 – – 20.82 ± 0.10 21.33 ± 0.12 20.45 ± 0.28 20.07 ± 0.44 19.59 ± 0.30 LT
20210903 59460.04 314.57 – – 21.10 ± 0.12 – 20.77 ± 0.15 20.48 ± 0.18 19.87 ± 0.16 Copernico 1.82-m
20210909 59466.07 320.42 – 22.05 ± 0.25 21.24 ± 0.25 21.62 ± 0.14 20.56 ± 0.20 20.55 ± 0.24 20.11 ± 0.27 TNG
20211027 59514.94 367.77 – – – 21.30 ± 0.06 20.55 ± 0.02 20.60 ± 0.07 19.95 ± 0.11 NOT
20220227 59637.92 486.94 – 22.33 ± 0.35 21.56 ± 0.25 21.85 ± 0.18 21.38 ± 0.18 21.18 ± 0.22 20.78 ± 0.15 NOT
Notes:
∗ Rest-frame phase in days from explosion, MJD=59135.42 ± 0.98.
𝐵𝑉 photometry is in the Vega system, while 𝑢𝑔𝑟𝑖𝑧 photometry is in the AB system.
Copernico 1.82-m: Copernico 1.82 m telescope at the Asiago Observatory (Italy). LT: 2-m Liverpool Telescope at the Observatorio del Roque de Los Muchachos (Spain). NOT: 2.56-m Nordic Optical Telescope
at the Roque de Los Muchachos Observatory (Spain). Schmidt Telescope: 67/91 Schmidt Telescope at the Asiago Observatory (Italy). TNG 3.6-m Telescopio Nazionale Galileo at the Roque de Los Muchachos
Observatory (Spain). TNT: 0.8-m Tsinghua University–NAOC (National Astronomical Observatories of China) Telescope at Xinglong Observatory of NAOC.
Table A2. 𝐽𝐻𝐾 Vega photometry of SN 2020wnt obtained with NOTCam.
UT date MJD Phase J H K
(days)∗ (mag) (mag) (mag)
20201216 59199.97 62.57 15.54 ± 0.01 15.57 ± 0.01 –
20210126 59240.36 101.71 15.37 ± 0.01 15.08 ± 0.01 14.965 ± 0.01
20210219 59264.43 125.03 16.23 ± 0.03 15.71 ± 0.03 15.577 ± 0.02
20210317 59290.39 150.18 16.26 ± 0.03 15.80 ± 0.03 15.747 ± 0.12
20210403 59307.40 166.67 16.64 ± 0.04 15.95 ± 0.11 15.960 ± 0.12
20210716 59411.40 267.44 17.49 ± 0.04 16.76 ± 0.05 16.340 ± 0.05
20210810 59436.31 291.58 18.39 ± 0.07 17.77 ± 0.06 17.275 ± 0.11
20211012 59499.15 352.47 19.16 ± 0.12 18.38 ± 0.15 17.808 ± 0.19
20211208 59556.30 407.85 19.56 ± 0.20 18.68 ± 0.20 18.083 ± 0.15
20220204 59614.40 464.15 19.77 ± 0.19 19.00 ± 0.18 17.991 ± 0.13
Notes:
∗ Rest-frame phase in days from explosion, MJD=59135.42 ± 0.98
Table A3. UV photometry obtained with Swift in the AB system.
UT date MJD Phase UVW2 UVM2 UVW1 U B V
(days)∗ (mag) (mag) (mag) (mag) (mag) (mag)
20201123 59176.18 39.51 22.12 ± 0.52 22.22 ± 0.56 19.95 ± 0.18 18.53 ± 0.13 17.82 ± 0.07 17.14 ± 0.09
20210114 59228.42 90.14 22.07 ± 0.50 > 21.64 21.14 ± 0.43 18.62 ± 0.13 17.53 ± 0.08 16.59 ± 0.08
20210203 59248.08 109.19 23.21 ± 1.51 > 21.43 – 20.58 ± 0.52 18.55 ± 0.09 17.16 ± 0.08
20210223 59268.20 128.68 22.47 ± 2.31 > 19.89 21.99 ± 3.16 20.78 ± 1.61 19.41 ± 0.30 18.07 ± 0.26
20210227 59272.37 132.73 22.94 ± 0.96 > 21.75 – – – –
20210303 59276.82 137.04 – > 21.39 21.69 ± 0.76 21.07 ± 0.80 19.51 ± 0.22 18.25 ± 0.24
20210315 59288.78 148.62 – > 21.45 – 22.37 ± 2.56 20.32 ± 0.63 18.06 ± 0.26
20210318 59291.03 150.80 22.11 ± 1.16 > 20.50 20.48 ± 0.57 20.65 ± 1.16 19.89 ± 0.23 18.29 ± 0.19
Notes:
∗ Rest-frame phase in days from explosion, MJD=59135.42 ± 0.98.
Table A4. ATLAS AB optical photometry.
UT date MJD Phase Band Magnitude
(days)∗ (mag)
20200926 59118.49 -16.386 𝑜 > 22.71
20201004 59126.48 -8.643 𝑜 > 22.00
20201014 59136.50 1.066 𝑐 19.70 ± 0.11
20201018 59140.48 4.922 𝑐 20.05 ± 0.22
20201020 59142.44 6.822 𝑜 19.61 ± 0.10
20201022 59144.41 8.731 𝑐 19.61 ± 0.09
20201024 59146.55 10.804 𝑜 19.02 ± 0.06
20201026 59148.50 12.694 𝑜 18.70 ± 0.11
20201026 59148.57 12.762 𝑐 19.70 ± 0.59
20201028 59150.59 14.719 𝑜 18.65 ± 0.04
20201031 59153.44 17.481 𝑜 18.52 ± 0.10
Table A4. – continued
20201109 59162.46 26.221 𝑜 17.81 ± 0.02
20201117 59170.48 33.992 𝑜 17.20 ± 0.02
20201123 59176.41 39.738 𝑜 16.92 ± 0.04
20201125 59178.48 41.744 𝑜 16.90 ± 0.01
20201127 59180.42 43.624 𝑜 16.86 ± 0.02
20201202 59185.47 48.517 𝑜 16.66 ± 0.02
20201203 59186.47 49.486 𝑜 16.66 ± 0.02
20201207 59190.35 53.246 𝑜 16.53 ± 0.01
20201209 59192.34 55.174 𝑐 16.79 ± 0.01
20201211 59194.41 57.180 𝑐 16.73 ± 0.01
20201217 59200.38 62.965 𝑐 16.67 ± 0.01
20201219 59202.54 65.058 𝑐 16.67 ± 0.01
20201223 59206.38 68.779 𝑜 16.33 ± 0.01
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 23
Table A4. – continued
20201224 59207.30 69.671 𝑜 16.30 ± 0.01
20201225 59208.36 70.698 𝑜 16.28 ± 0.01
20201231 59214.35 76.502 𝑜 16.27 ± 0.01
20210102 59216.33 78.421 𝑜 16.28 ± 0.01
20210104 59218.39 80.417 𝑜 16.28 ± 0.01
20210105 59219.39 81.386 𝑐 16.64 ± 0.01
20210106 59220.32 82.287 𝑐 16.63 ± 0.01
20210108 59222.31 84.215 𝑐 16.72 ± 0.01
20210112 59226.35 88.130 𝑐 16.78 ± 0.01
20210114 59228.39 90.107 𝑐 16.83 ± 0.01
20210116 59230.30 91.957 𝑐 16.86 ± 0.01
20210128 59242.37 103.653 𝑜 16.54 ± 0.03
20210130 59244.34 105.562 𝑜 16.55 ± 0.01
20210131 59245.26 106.453 𝑜 16.55 ± 0.01
20210201 59246.29 107.452 𝑜 16.58 ± 0.01
20210205 59250.30 111.337 𝑐 17.42 ± 0.01
20210207 59252.32 113.295 𝑐 17.56 ± 0.02
20210209 59254.29 115.203 𝑐 17.65 ± 0.01
20210211 59256.27 117.122 𝑐 17.83 ± 0.03
20210213 59258.30 119.089 𝑐 17.88 ± 0.02
20210217 59262.26 122.926 𝑜 17.31 ± 0.02
20210223 59268.30 128.779 𝑜 17.54 ± 0.04
20210227 59272.24 132.597 𝑜 17.60 ± 0.03
20210303 59276.24 136.473 𝑜 17.60 ± 0.03
20210311 59284.26 144.244 𝑐 18.41 ± 0.04
20210325 59298.23 157.781 𝑜 17.78 ± 0.05
20210327 59300.25 159.738 𝑜 17.81 ± 0.06
20210711 59406.60 262.791 𝑜 18.73 ± 0.09
20210719 59414.61 270.552 𝑜 18.72 ± 0.07
20210721 59416.60 272.481 𝑜 19.12 ± 0.28
20210727 59422.61 278.304 𝑜 19.69 ± 0.23
20210806 59432.61 287.994 𝑜 20.07 ± 0.22
20210808 59434.62 289.942 𝑐 20.46 ± 0.19
20210812 59438.62 293.818 𝑐 20.82 ± 0.31
20210820 59446.61 301.560 𝑜 20.76 ± 0.30
20210905 59462.56 317.016 𝑜 21.09 ± 0.35
20210911 59468.58 322.849 𝑐 21.05 ± 0.38
20210915 59472.56 326.705 𝑐 21.71 ± 0.71
20211011 59498.50 351.841 𝑜 20.42 ± 0.30
20211015 59502.47 355.688 𝑜 20.18 ± 0.34
20211017 59504.43 357.587 𝑜 19.99 ± 0.29
20211104 59522.52 375.116 𝑜 20.17 ± 0.28
20211106 59524.47 377.006 𝑐 21.64 ± 0.65
20220204 59614.28 464.031 𝑐 22.00 ± 0.97
20220205 59615.28 465.000 𝑐 21.96 ± 1.07
Notes:
∗ Rest-frame phase in days from explosion, MJD=59135.42 ± 0.98.
Table A5. ZTF AB optical photometry.
UT date MJD Phase Band Magnitude
(days)∗ (mag)
20201014 59136.40 0.95 𝑔 19.63 ± 0.06
20201014 59136.47 1.02 𝑟 19.57 ± 0.05
20201017 59139.37 3.83 𝑟 19.96 ± 0.09
20201017 59139.41 3.87 𝑔 20.24 ± 0.13
20201019 59141.34 5.74 𝑟 19.90 ± 0.08
20201019 59141.42 5.81 𝑔 20.29 ± 0.09
20201021 59143.32 7.66 𝑟 19.62 ± 0.06
20201021 59143.39 7.72 𝑔 19.90 ± 0.06
20201023 59145.34 9.61 𝑟 19.21 ± 0.03
20201023 59145.42 9.69 𝑔 19.65 ± 0.05
Table A5. – continued
UT date MJD Phase Band Magnitude
(days)∗ (mag)
20201027 59149.37 13.52 𝑟 18.83 ± 0.05
20201027 59149.46 13.60 𝑔 19.36 ± 0.06
20201029 59151.33 15.42 𝑔 19.02 ± 0.08
20201029 59151.36 15.45 𝑟 18.67 ± 0.05
20201031 59153.34 17.36 𝑟 18.58 ± 0.05
20201031 59153.40 17.42 𝑔 19.03 ± 0.11
20201102 59155.29 19.25 𝑟 18.39 ± 0.04
20201102 59155.39 19.35 𝑔 19.00 ± 0.09
20201104 59157.39 21.29 𝑔 18.92 ± 0.08
20201104 59157.44 21.34 𝑟 18.29 ± 0.04
20201111 59164.32 28.00 𝑟 17.70 ± 0.01
20201111 59164.43 28.11 𝑔 18.17 ± 0.02
20201113 59166.30 29.92 𝑟 17.77 ± 0.09
20201113 59166.41 30.03 𝑔 17.99 ± 0.02
20201115 59168.26 31.82 𝑔 17.89 ± 0.02
20201115 59168.30 31.86 𝑟 17.42 ± 0.01
20201117 59170.23 33.73 𝑔 17.73 ± 0.01
20201117 59170.28 33.78 𝑟 17.30 ± 0.01
20201119 59172.23 35.67 𝑔 17.63 ± 0.01
20201122 59175.30 38.64 𝑔 17.42 ± 0.02
20201122 59175.37 38.71 𝑟 17.00 ± 0.01
20201124 59177.39 40.67 𝑔 17.38 ± 0.01
20201126 59179.25 42.47 𝑔 17.33 ± 0.02
20201126 59179.31 42.53 𝑟 16.95 ± 0.01
20201128 59181.22 44.38 𝑔 17.35 ± 0.03
20201128 59181.28 44.44 𝑟 16.88 ± 0.02
20201130 59183.21 46.31 𝑔 17.12 ± 0.03
20201202 59185.19 48.23 𝑔 17.08 ± 0.02
20201202 59185.26 48.29 𝑟 16.70 ± 0.01
20201204 59187.20 50.17 𝑔 17.10 ± 0.01
20201204 59187.24 50.21 𝑟 16.72 ± 0.01
20201206 59189.22 52.13 𝑟 16.60 ± 0.01
20201209 59192.16 54.98 𝑟 16.64 ± 0.02
20201209 59192.22 55.04 𝑔 17.06 ± 0.02
20201211 59194.22 56.98 𝑔 16.98 ± 0.01
20201211 59194.26 57.02 𝑟 16.53 ± 0.01
20201214 59197.24 59.90 𝑟 16.44 ± 0.02
20201217 59200.18 62.75 𝑔 16.79 ± 0.01
20201217 59200.24 62.81 𝑟 16.40 ± 0.01
20201220 59203.22 65.70 𝑟 16.45 ± 0.01
20201220 59203.31 65.78 𝑔 16.88 ± 0.01
20201222 59205.15 67.57 𝑟 16.38 ± 0.01
20201222 59205.26 67.67 𝑔 16.79 ± 0.01
20201224 59207.18 69.53 𝑔 16.95 ± 0.16
20201224 59207.20 69.55 𝑟 16.55 ± 0.13
20201228 59211.16 73.39 𝑟 16.38 ± 0.01
20201228 59211.22 73.45 𝑔 16.81 ± 0.02
20210102 59216.18 78.26 𝑔 16.69 ± 0.01
20210102 59216.28 78.35 𝑟 16.30 ± 0.01
20210104 59218.23 80.24 𝑔 16.81 ± 0.01
20210104 59218.35 80.36 𝑟 16.32 ± 0.01
20210106 59220.20 82.15 𝑔 16.82 ± 0.01
20210106 59220.24 82.19 𝑟 16.41 ± 0.01
20210108 59222.18 84.07 𝑔 16.89 ± 0.02
20210108 59222.20 84.09 𝑟 16.35 ± 0.01
20210110 59224.20 86.03 𝑔 17.05 ± 0.01
20210110 59224.25 86.08 𝑟 16.44 ± 0.01
20210112 59226.21 87.97 𝑔 17.07 ± 0.01
20210112 59226.24 88.00 𝑟 16.48 ± 0.01
20210114 59228.20 89.90 𝑔 17.10 ± 0.01
20210114 59228.22 89.92 𝑟 16.49 ± 0.01
20210116 59230.20 91.84 𝑔 17.17 ± 0.01
20210116 59230.22 91.86 𝑟 16.50 ± 0.01
MNRAS 000, 1–24 (2022)
24 Gutiérrez et al.
Table A5. – continued
UT date MJD Phase Band Magnitude
(days)∗ (mag)
20210118 59232.18 93.76 𝑟 16.48 ± 0.01
20210118 59232.22 93.80 𝑔 17.26 ± 0.01
20210129 59243.24 104.48 𝑔 17.53 ± 0.18
20210204 59249.14 110.19 𝑔 18.06 ± 0.07
20210204 59249.19 110.24 𝑟 16.90 ± 0.01
20210206 59251.14 112.13 𝑔 18.10 ± 0.02
20210206 59251.18 112.17 𝑟 16.95 ± 0.01
20210208 59253.16 114.09 𝑟 17.04 ± 0.01
20210208 59253.18 114.11 𝑔 18.22 ± 0.02
20210210 59255.14 116.01 𝑟 17.14 ± 0.01
20210210 59255.16 116.03 𝑔 18.40 ± 0.02
20210212 59257.14 117.95 𝑟 17.28 ± 0.01
20210212 59257.16 117.97 𝑔 18.52 ± 0.02
20210215 59260.18 120.89 𝑔 18.68 ± 0.04
20210215 59260.22 120.93 𝑟 17.35 ± 0.02
20210218 59263.16 123.78 𝑔 18.81 ± 0.07
20210218 59263.17 123.79 𝑟 17.53 ± 0.02
20210220 59265.14 125.70 𝑟 17.57 ± 0.01
20210220 59265.18 125.74 𝑔 18.92 ± 0.05
20210222 59267.14 127.64 𝑟 17.64 ± 0.02
20210224 59269.14 129.57 𝑟 17.62 ± 0.02
20210224 59269.25 129.68 𝑔 19.02 ± 0.10
20210226 59271.14 131.51 𝑟 17.72 ± 0.02
20210228 59273.16 133.47 𝑔 19.24 ± 0.12
20210228 59273.21 133.52 𝑟 17.80 ± 0.04
20210302 59275.16 135.41 𝑔 19.35 ± 0.09
20210302 59275.26 135.50 𝑟 17.78 ± 0.04
20210305 59278.14 138.29 𝑟 17.86 ± 0.03
20210305 59278.21 138.36 𝑔 19.34 ± 0.11
20210318 59291.19 150.94 𝑟 17.91 ± 0.02
20210320 59293.18 152.87 𝑟 17.91 ± 0.03
20210323 59296.20 155.79 𝑔 19.37 ± 0.20
20210325 59298.16 157.69 𝑟 17.96 ± 0.08
20210625 59390.48 247.15 𝑟 18.89 ± 0.14
20210626 59391.48 248.12 𝑟 19.04 ± 0.17
20210627 59392.48 249.09 𝑟 18.92 ± 0.16
21 Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer
de Can Magrans s/n, 08193 Barcelona, Spain
22 School of Physics, O’Brien Centre for Science North, University
College Dublin, Dublin, Ireland.
23 Department of Physics, Florida State University, 77 Chieftan
Way, Tallahassee, FL 32306, USA
24 Turku Collegium for Science, Medicine and Technology,
University of Turku, FI-20014 Turku, Finland
25 National Astronomical Observatories, Chinese Academy of
Sciences, Beijing 100101, China
26 School of Astronomy and Space Science, University of Chinese
Academy of Sciences, Beijing 101408, China
27 Beijing Planetarium, Beijing Academy of Science and Technol-
ogy, Beijing 100044, China
28 Yunnan Observatories, Chinese Academy of Sciences, Kunming
650216, China
29 Key Laboratory for the Structure and Evolution of Celestial
Objects, Chinese Academy of Sciences, Kunming 650216, China
This paper has been typeset from a TEX/LATEX file prepared by the author.
Table A5. – continued
UT date MJD Phase Band Magnitude
(days)∗ (mag)
20210628 59393.48 250.06 𝑟 18.91 ± 0.10
20210629 59394.48 251.03 𝑟 18.85 ± 0.09
20210630 59395.48 252.00 𝑟 18.94 ± 0.10
20210701 59396.48 252.97 𝑟 18.89 ± 0.09
20210702 59397.48 253.93 𝑟 18.94 ± 0.09
20210704 59399.48 255.87 𝑟 19.02 ± 0.10
20210705 59400.48 256.84 𝑟 19.02 ± 0.12
20210706 59401.48 257.81 𝑟 18.92 ± 0.09
20210709 59404.47 260.71 𝑟 18.90 ± 0.08
20210710 59405.47 261.68 𝑟 19.02 ± 0.09
20210711 59406.47 262.65 𝑟 18.90 ± 0.06
20210731 59426.45 282.01 𝑟 19.64 ± 0.10
20210808 59434.47 289.78 𝑔 20.66 ± 0.14
20210808 59434.47 289.78 𝑟 20.27 ± 0.12
20210808 59434.47 289.78 𝑟 20.27 ± 0.12
20210810 59436.47 291.72 𝑟 20.40 ± 0.10
20210810 59436.47 291.72 𝑟 20.40 ± 0.10
20210817 59443.48 298.51 𝑟 20.79 ± 0.19
20210817 59443.48 298.51 𝑟 20.79 ± 0.19
20210930 59487.39 341.06 𝑟 20.67 ± 0.20
20210930 59487.39 341.06 𝑟 20.67 ± 0.20
20211009 59496.41 349.80 𝑟 20.43 ± 0.17
20211009 59496.41 349.80 𝑟 20.43 ± 0.17
20211011 59498.37 351.70 𝑟 20.68 ± 0.13
20211011 59498.37 351.70 𝑟 20.68 ± 0.13
20211016 59503.40 356.57 𝑟 20.65 ± 0.20
20211016 59503.40 356.57 𝑟 20.65 ± 0.20
20211030 59517.34 370.08 𝑟 21.21 ± 0.18
20211030 59517.34 370.08 𝑟 21.21 ± 0.18
20211103 59521.29 373.91 𝑟 20.92 ± 0.15
20211103 59521.29 373.91 𝑟 20.92 ± 0.15
20211107 59525.31 377.80 𝑟 21.04 ± 0.17
20211107 59525.31 377.80 𝑟 21.04 ± 0.17
Notes:
∗ Rest-frame phase in days from explosion, MJD=59135.42 ± 0.98.
MNRAS 000, 1–24 (2022)
SN 2020wnt: an unusual SLSN-I 25
Table A6. Spectroscopic observations of SN 2020wnt
UT date MJD Rest-frame Range Telescope Grism/Grating
phase (days)∗ (Å) +Instrument
20201115 59168.00 31.6 3205 − 9775 LICK+KAST –
20201121 59174.92 38.3 3390 − 7915 Ekar+AFOSC gm4
20201124 59177.13 40.4 3550 − 9005 Ekar+AFOSC VPH7
20201129 59182.79 45.9 3680 − 7945 Ekar+AFOSC gm4
20201202 59185.56 48.6 6700 − 24620 IRTF+Spex Prism
20201203 59186.10 49.1 3430 − 10075 TNG+LRS LR-B+LR-R
20201212 59195.10 57.8 3295 − 9335 NOT+ALFOSC Grism#4
20201215 59198.05 60.7 3440 − 9350 NOT+ALFOSC Grism#4
20201216 59199.85 62.4 4845 − 9010 Ekar+AFOSC VPH6
20201218 59201.10 63.6 2975 − 10070 TNG+LRS LR-B+LR-R
20201220 59203.78 66.2 4845 − 9010 Ekar+AFOSC VPH6
20210101 59215.70 77.8 3390 − 8495 LJT+YFOSC G3
20210102 59216.44 78.5 3740 − 8550 XLT+BFOSC G4
20210110 59224.92 86.7 3250 − 9010 Ekar+AFOSC VPH6
20210114 59228.53 90.2 3390 − 8490 LJT+YFOSC G3
20210115 59229.48 91.1 3740 − 8540 XLT+BFOSC G4
20210118 59232.76 94.3 3295 − 9010 Ekar+AFOSC VPH6
20210122 59236.53 98.0 3390 − 8495 LJT+YFOSC G3
20210127 59241.89 103.2 3295 − 9375 NOT+ALFOSC Grism#4
20210204 59249.69 110.7 3390 − 8490 LJT+YFOSC G3
20210211 59256.52 117.3 3390 − 8495 LJT+YFOSC G3
20210216 59261.84 122.5 3295 − 9350 NOT+ALFOSC Grism#4
20210311 59284.86 144.8 3295 − 9385 NOT+ALFOSC Grism#4
20210317 59290.86 150.6 3000 − 10050 TNG+LRS LR-B+LR-R
20210407 59311.80 170.9 4845 − 9010 Ekar+AFOSC VPH6
20210728 59423.20 278.9 3540 − 7630 GTC + OSIRIS R1000B
20210811 59438.16 293.3 4945 − 10100 GTC + OSIRIS R1000R
NOTES:
∗ Rest-frame phase in days from explosion, MJD=59135.42 ± 0.98.
Telescope code: EKAR: Copernico 1.82-m Telescope, INAF (Mount Ekar); AFOSC: Asiago Faint Object Spectrograph and Camera; IRTF: NASA Infrared
Telescope Facility; SpeX: 0.7-5.3 Micron Medium-Resolution Spectrograph and Imager; LICK: 3-m Shane telescope at Lick Observatory; KAST: Kast
double spectrograph; LJT: Lijiang 2.4-m Telescope; YFOSC: Yunnan Faint Object Spectrograph and Camera; XLT: Xinglong 2.16-m Telescope; BFOSC:
Beijing Faint Object Spectrograph and Camera; NOT: Nordic Optical Telescope; ALFOSC: Alhambra Faint Object Spectrograph and Camera; TNG: 3.6-
m Telescopio Nazionale Galileo, INAF; LRS: Low Resolution Spectrograph; GTC: Gran Telescopio Canarias; OSIRIS: Optical System for Imaging and
low-Intermediate-Resolution Integrated Spectroscopy.
MNRAS 000, 1–24 (2022)