# Probing the Circumstellar Environment of highly luminous type IIn SN
ASASSN-14il

Naveen Dukiya Aryabhatta Research Institute of Observational Sciences, Manora
Peak 263001, India  Department of Applied Physics, Mahatma Jyotiba Phule
Rohilkhand University, Bareilly, 243006, India  Anjasha Gangopadhyay Hiroshima
Astrophysical Science Centre, Hiroshima University, 1-3-1 Kagamiyama, Higashi-
Hiroshima, Hiroshima 739-8526, Japan  Kuntal Misra Aryabhatta Research
Institute of Observational Sciences, Manora Peak 263001, India  Griffin
Hosseinzadeh Steward Observatory, University of Arizona, 933 North Cherry
Avenue, Tucson, AZ 85721-0065, USA  K. Azalee Bostroem Steward Observatory,
University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA
Bhavya Ailawadhi Aryabhatta Research Institute of Observational Sciences,
Manora Peak 263001, India  Department of Physics, Deen Dayal Upadhyaya
Gorakhpur University, Gorakhpur, 273009, India  D. Andrew Howell Las Cumbres
Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA
Department of Physics, University of California, Santa Barbara, CA 93106-9530,
USA  Stefano Valenti Department of Physics and Astronomy, University of
California, Davis, 1 Shields Avenue, Davis, CA 95616-5270, USA  Iair Arcavi
School of Physics and Astronomy, Tel Aviv University, Tel Aviv 69978, Israel
Curtis McCully Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta,
CA 93117-5575, USA  Department of Physics, University of California, Santa
Barbara, CA 93106-9530, USA

(Revised –; Accepted –)

###### Abstract

We present long-term photometric and spectroscopic studies of Circumstellar
Material (CSM)-Ejecta interacting supernova (SN) ASASSN-14il in the galaxy PGC
3093694. The SN reaches a peak $r$-band magnitude of $\sim$ $-20.3\pm 0.2$ mag
rivaling SN 2006tf and SN 2010jl. The multiband and the pseudo-bolometric
lightcurve show a plateau lasting $\sim 50$ days. Semi-analytical CSM
interaction models can match the high luminosity and decline rates of the
lightcurves but fail to faithfully represent the plateau region and the bumps
in the lightcurves. The spectral evolution resembles the typical SNe IIn
dominated by CSM interaction, showing blue-continuum and narrow Balmer lines.
The lines are dominated by electron scattering at early epochs. The signatures
of the underlying ejecta are visible as the broad component in the H$\alpha$
profile from as early as day 50, hinting at asymmetry in the CSM. A narrow
component is persistent throughout the evolution. The SN shows remarkable
photometric and spectroscopic similarity with SN 2015da. However, the
different polarization in ASASSN-14il compared to SN 2015da suggests an
alternative viewing angle. The late-time blueshift in the H$\alpha$ profiles
supports dust formation in the post-shock CSM or ejecta. The mass-loss rate of
2-7 M${}_{\odot}\mathrm{yr}^{-1}$ suggests a Luminous Blue Variable (LBV)
progenitor in an eruptive phase for ASASSN-14il.

supernovae: general – supernovae: individual: ASASSN-14il – galaxies:
individual: – techniques: photometric – techniques: spectroscopic

##  1 Introduction

Core-collapse supernovae (CCSNe) that show narrow (10s to 1000s of km/s)
Balmer lines in their spectra are typically classified as type IIn supernovae
(hereafter SNe IIn; Schlegel, 1990; Filippenko, 1997). The narrow lines in
their spectra arise due to the ejecta of the SN interacting with dense
Circumstellar-Material (CSM) surrounding the progenitor (Chugai & Danziger,
1994; Chugai, 2001; Chugai et al., 2004). Early spectra are dominated by a
blue continuum, gradually cooling down as the SN evolves. The line profiles
are usually multi-component and asymmetric based on the geometry of the
explosion. The narrow component (10s to 100s km sec-1) results from the
recombination of ionized unshocked CSM ahead of the blast-wave shock. This
usually carries a P-Cygni profile and characterizes the velocity of CSM wind.
The intermediate width component (1000s of km sec-1) arises from the radiation
emitted by the Cold-Dense Shell (CDS) region between the forward and the
reverse shock and characterizes the velocity at which the blast wave is moving
through the CSM (Chevalier & Fransson, 1994). Sometimes, broad components from
the underlying explosion, i.e., ejecta signatures, are visible due to
diminished interaction or asymmetric CSM geometry (see Smith, 2017 for a
review on SNe IIn). These interacting SNe also radiate in X-ray and radio
bands in addition to Ultraviolet, Optical, and Infrared. However, the
radiation at these wavelengths can be self-absorbed and reprocessed,
especially at early times owing to the higher density of CSM (Chevalier, 1982,
1998). Radio and X-ray detections of these events, however, can provide
independent constraints on the mass loss of the progenitor. Some SNe IIn also
form dust at the later stages of evolution, which leads to IR excess at late
times or dust-echoes (Pozzo et al., 2004; Mattila et al., 2008; Fransson et
al., 2014; Tartaglia et al., 2020).

SNe IIn are relatively rare ($\sim$9% intrinsic rate among all SNe II; Li et
al., 2011) and form a heterogeneous and poorly understood class of transients
due to their exceptionally diverse properties. These events have typical
absolute magnitudes ranging from M${}_{\mathrm{R}}\sim$ $-16$ to $-19$ mag but
some unusually luminous events even reach M${}_{\mathrm{R}}\sim$ $-22$ mag
(Kiewe et al., 2012; Taddia et al., 2013; Smith & McCray, 2007). In addition,
the overall duration and shape of the lightcurve vary a lot. Some SNe IIn are
extremely bright with long-lasting lightcurves (SN 2006gy - Smith et al.,
2007; SN 2006tf - Smith et al., 2008); some are slow evolving like (SN 1988Z -
Chugai & Danziger, 1994; SNe 2005ip and 2006jd - Stritzinger et al., 2012).
Some SNe IIn shows linear decline (SN 1999el - Di Carlo et al., 2002) and some
display a plateau (SN 1994W - Dessart et al., 2009; SN 2009kn - Kankare et
al., 2012; 2011ht - Mauerhan et al., 2013). Depending on their lightcurve
shapes and luminosities, different physical scenarios have been invoked in the
literature to explain the powering mechanism of these SNe IIn. A lightcurve
powered by ejecta interacting with optically thin CSM can explain moderately
bright SNe IIn, but superluminous SNe IIn events require extreme conditions.
In order to explain the extreme brightness of SN 2006gy, Smith & McCray (2007)
proposed the interaction of ejecta with a compact dense shell as a powering
mechanism where the photons diffuse through the optically thick CSM. Woosley
et al. (2007) argues that the power source is the collision between shells
ejected by a pulsational pair-instability SNe instead. Both these cases point
to a very massive progenitor and dense CSM shells. If the CSM shells are
extended, they can give rise to the high-luminosities and a relatively slower
decline (as seen in SNe 2006tf and 2010jl; Smith et al., 2008; Fransson et
al., 2014). Recently, Nicholl et al. (2020) and Suzuki et al. (2021) showed
that the lightcurve of SN 2016aps requires a CSM mass $>$ 40M⊙ to power its
luminosity. The dense CSM from the progenitors’ mass-loss events may not
extend far. In that case, a transition from an interaction-dominated regime to
a Ni-powered regime is seen as similar to SNe IIn-P (Kankare et al., 2012;
Mauerhan et al., 2013). However, all these events share the fact that a
significant part of the total radiated energy during its lifetime comes from
the CSM interaction of ejecta.

Different kinds of progenitor systems can give rise to SNe IIn. They have
different CSM profiles (created by the mass-loss events of the progenitor or
its binary companion) resulting from these different progenitor systems. SNe
IIn typically requires mass-loss rates higher than
$10^{-3}\mathrm{M}_{\odot}\mathrm{yr}^{-1}$ (Moriya et al., 2014) which is
higher than what can be expected from the line-driven winds (Smith & Owocki,
2006), although SNe IIn falling in the lower range of luminosity can be
explained by long-term strong winds (Fransson et al., 2002; Smith et al.,
2009a). The enhanced mass-loss from Luminous Blue Variable (LBV) eruptions in
months to decades before an SN explosion (Smith, 2014) make them much more
suitable progenitor candidates for most SNe IIn. In the case of many SNe,
archival images from Hubble Space Telescope (HST) and pre-SN outbursts have
revealed their progenitors to be consistent with LBVs (Gal-Yam & Leonard,
2009; Smith et al., 2010, 2011; Kochanek et al., 2011). Additionally, the
large amount of CSM required to explain the superluminous supernovae (SLSNe)
also favors a massive (40 - 100 M⊙) LBV progenitor. The outburst in LBVs is
still not well understood, but for the more massive LBVs (70-140 M⊙), it can
be due to pulsational-pair instability (Woosley, 2017; Vigna-Gómez et al.,
2019). Alternatively, the CSM can result from binary interactions and common
envelope ejection. This is usually seen in Ia-CSM type of SNe where a
thermonuclear SN explodes within the CSM generated by its companion (SN 2002ic
- Benetti et al., 2006; iPTF11kx - Silverman et al., 2013; SNe 2012ca and
2013dn - Fox et al., 2015).

Many authors have developed analytical and numerical modeling frameworks to
explain the explosion physics and observables of SNe IIn (Balberg & Loeb,
2011; Svirski et al., 2012; Chatzopoulos et al., 2012; Moriya et al., 2013;
Ofek et al., 2014; Dessart et al., 2015; Jiang et al., 2020). Continued
observations during the evolution of such events can decipher the mass-loss
history decades before the SN explosion, giving crucial insight into the final
stages of stellar evolution. In some sense, these interacting SNe are the only
real probes to the pre-SN physics of their exotic progenitors. This is
especially true at higher redshifts where identifying progenitors using direct
imaging is difficult.

In this context, we present the long-term observational photometric and
spectroscopic analysis of ASASSN-14il. It was discovered by All Sky Automated
Survey for SuperNovae (ASAS-SN) at UT 2014-10-01.11 at $V\sim 16.5$ mag
(Brimacombe et al., 2014). The transient was detected in images from the
double 14-cm “Cassius” telescope in Cerro Tololo, Chile at RA = 00h45m32.55s,
Dec = -14d15m34.6s; approximately 0′′.33 North and 0′′.26 West from the centre
of the galaxy 2MASX J00453260-1415328. No source was detected at the position
of the transient down to the limiting magnitude of $V\sim 17.3$ mag in images
taken on UT 2014-09-27.04 and before. Spectroscopic observation on UT
2014-10-03.55 with the Wide Field Spectrograph (WiFeS) mounted on the
Australian National University 2.3-m telescope, using the B3000/R3000 gratings
($3500-9800$ Å, $1$ Å resolution) was used to classify the SN (Childress et
al., 2014). ASASSN-14il was classified as a SN IIn based on the narrow
emission lines ($\sim$3000 km/s) in the Balmer series and the blue continuum.
The redshift calculated from the SN spectrum is consistent with the redshift
of the host at 0.022 (Jones et al., 2009). The redshift of 0.022 corresponds
to a distance of $88.5\pm 6.2$ Mpc111NED(corrected for Virgo + GA + Shapley)
assuming H${}_{0}=73$ km/s/Mpc, $\Omega_{\mathrm{matter}}=0.27$, and
$\Omega_{\mathrm{vaccum}}=0.73$. This distance is adopted throughout the paper
for further analysis.

The paper is structured as follows. Section 2 describes the observations and
data reduction procedures. Estimation of extinction and explosion epoch is
discussed in section 3. The lightcurve evolution and modeling in described in
sections 4 and 5, respectively. The spectral evolution is described in section
6. Finally, we conclude our paper in section 7 and summarize the results in
section 8.

##  2 Observations and Data Reduction

Table 1: General information about ASASSN-14il and its host galaxy

ASASSN-14il
---
Discovery Date(1) | UT 2014-10-01.11
Explosion Date | UT 2014-09-25.9
SN Type(2) | IIn
RA (J2000) | 00h45m32.55s
Dec (J2000) | -14d15m34.6s
Discovery Magnitude | 16.5 mag (V-band)
E(B-V)total | 0.21 ± 0.08 mag
Host Galaxy(3)
Galaxy Names | PGC 3093694
| WISEA J004532.52-141532.7
| 6dF J0045326-141533
Morphology Type | Dwarf Galaxy
RA (J2000) | 00h45m32.601s
Dec (J2000) | -14d15m32.50s
redshift (z) | 0.022
vVirgo+Shapley+GA | 6462 ± 47 km/s
D |  $88.52\pm 6.2$ Mpc
$\mu$ | 34.74 ± 0.15 mag

References - (1) Brimacombe et al. (2014), (2) Childress et al. (2014), (3)
NED

We observed ASASSN-14il with the Las Cumbres Observatory (LCO) network of
telescopes as part of the Supernova Key Project (which eventually became the
Global Supernova Project). Observations started from $\sim 2$ days after
discovery and continued for over a year with a high cadence in BVgri filters.
Since the source was contaminated by the host galaxy, we performed image
subtraction using High Order Transform of PSF ANd Template Subtraction
(HOTPANTS; Becker, 2015) algorithm integrated in the
lcogtsnpipe222https://github.com/LCOGT/lcogtsnpipe/ pipeline (Valenti et al.,
2016). The template images were taken in December 2022, long after the SN
faded. Figure 1 shows a gri composite image of the science frame, reference
frame, and difference image of ASASSN-14il. The BVgri instrumental magnitudes
were obtained from the difference images and were calibrated against the AAVSO
Photometric All-Sky Survey (APASS) catalog. The photometry of ASASSN-14il is
presented in Table 7.

Figure 1: The left panel shows the gri color-composite image of ASASSN-14il
with frames taken on 28 November 2014. The central panel shows the gri color-
composite with the reference image taken on 27 December 2022, and the right
panel shows the color-composite of the difference images with the position of
SN marked. Figure 2: The spectrum of ASASSN-14il showing a prominent dip at
5890 Å and 5896 Å (marked as orange) in the rest-frame of the galaxy
indicating high host galaxy extinction. The rest-frame wavelength of the
nearby HeI emission line is marked in blue. The grey line shows the continuum
used for normalization.

This supernova was also observed using the Ultra-Violet/Optical Telescope
(UVOT; Roming et al., 2005) on The Neil Gehrels Swift Observatory (hereafter
Swift; Gehrels et al. 2004) in UVW2, UVM2, UVW1, U, B, and $V$ bands starting
from day $\sim 8$ to $49$ after the explosion (c.f.r section 3). The
photometry was obtained from the Swift Optical/Ultraviolet Supernova Archive
(SOUSA; Brown et al., 2014). The reduction is based on that of Brown et al.
(2009), which includes subtraction of the host galaxy count rates (from
templates taken $\sim$ 2 years after the explosion) using a 5″aperture at the
source position, using the revised UV zero points and time-dependent
sensitivity from Breeveld et al. (2011). Aperture photometry was performed
using an aperture size of 3 or 5 arcsec based on the error. The final UVOT
magnitudes are given in Table 8.

Low-resolution (R $\sim$ 400-700) optical spectroscopic observations were
carried out using the FLOYDS spectrographs mounted on the LCO 2m telescopes.
The 1D wavelength and flux calibrated spectra were extracted using the
floydsspec333https://github.com/LCOGT/floyds_pipeline pipeline (Valenti et
al., 2014). We include the publicly available data taken under the Public ESO
Spectroscopic Survey of Transient Objects (PESSTO) program (Smartt et al.,
2015) using EFOSC2 mounted on the 3.6m ESO-NTT. We also include the high
resolution (R $\sim$ 3000) spectroscopic data, available in WISeREP (Yaron &
Gal-Yam, 2012), which were taken with the Wide Field Spectrograph on the
Australian National University 2.3-m telescope under the ANU WiFeS SuperNovA
Programme (AWSNAP; Childress et al., 2016). Our optical spectroscopic
observations span from day $\sim 8$ to $327$ after the explosion. NIR
spectroscopy was acquired using SOFI mounted on the 3.6m ESO-New Technology
Telescope (NTT) under the PESSTO program. The NIR observations span from day
$\sim 19$ to $\sim 116$. The log of spectroscopic observations is given in
Table 9. All the spectra were scaled to photometry by using lightcurve-
fitting444https://github.com/griffin-h/lightcurve_fitting module (Hosseinzadeh
& Gomez, 2022) to account for the slit loss corrections. Finally, all the
spectra were corrected for the heliocentric redshift of the host galaxy.

##  3 Estimation of extinction and explosion epoch

The Milky Way extinction along the line-of-sight of ASASSN-14il is
E(B-V)${}_{\mathrm{MW}}=0.0019$ (Schlafly & Finkbeiner, 2011). To estimate the
host galaxy contribution to the total reddening, we search for the NaID 5890,
5896 Å doublet in the high-resolution ANU-WiFeS spectra. The NaID doublet is
seen at the corresponding host-galaxy rest wavelength. We combine the ANU
spectra taken on UT 2014-10-09.50, UT 2014-10-18.56, and UT 2014-10-27.47 to
increase the signal-to-noise ratio. The NaID doublet is contaminated by the
broad wings of a nearby emission line, as seen in figure Figure 2.
Additionally, we note that the HeI emission is blueshifted from its rest
wavelength, the implications of which is discussed in section 6. We model the
wings of the emission line as a part of the continuum and estimate the
equivalent width (EW) of Na D1 and D2 from the normalized spectra to be
$0.44\pm 0.03$ Å, and $0.57\pm 0.03$ Å, respectively. We use the relation
given by Poznanski et al. (2012) between the Na1D features and $E(B-V)$. From
this relation, we estimate host galaxy reddening of $0.22\pm 0.09$ mag and
$0.21\pm 0.08$ mag, respectively, for D1 and D2. We also use the relation for
the combined equivalent width of D1 and D2 to get the reddening value of
$0.22\pm 0.09$ mag, similar to the one estimated individually from D1 and D2.
We multiply this reddening value by 0.86 to be consistent with the
recalibration of Milkyway extinction by Schlafly & Finkbeiner (2011).
Therefore, the host-galaxy extinction is $0.19\pm 0.08$ mag. The fitting
errors have been propagated in quadrature. Thus, we adopt a total
(galaxy+host) $E(B-V)=0.21\pm 0.08$ mag. We use this value of extinction
throughout the paper.

To estimate the explosion epoch of ASASSN-14il, we perform parabolic fitting
on the early time $V$-band lightcurve. The magnitudes are converted to fluxes
up to 25 days post-discovery. The best-fit coefficients are used to find the
roots of the equation, i.e., the value of time for which the flux equals zero.
We perform the fit using MCMC simulations to constrain the associated errors.
The explosion epoch is estimated to be JD $2456926.4\pm 0.5$ (UT 2014-09-25.9)
and adopted as the reference epoch throughout the paper. The general
information about ASASSN-14il and its host galaxy is listed in Table 1.

##  4 Lightcurve Characteristics of asassn-14il

Figure 3: The evolution of optical and UV lightcurves of ASASSN-14il. Different bands are plotted with an offset for clarity. Circles represent the photometry from LCOGT, and squares represent the photometry from Swift UVOT. Epochs having accompanying spectra are represented by line symbols at the bottom. A plateau can be seen in the optical bands between days 40–80. A prominent bump in the lightcurve can be seen around day 90. Table 2: The rise time and peak magnitudes of UV-Optical bands with a clear peak observed. The listed value and error of Rise Time and Observed Peak Magnitude are derived from the distribution resulting from the MC experiments. | Rise Time | Peak Magnitude
---|---|---
|  | Observed | Absolute
| (day) | (mag) | (mag)
UVW2 | $14.7\pm 0.9$ | $14.54\pm 0.03$ | $-22.0\pm 0.7$
UVM2 | $15.2\pm 0.9$ | $14.43\pm 0.03$ | $-22.2\pm 0.7$
UVW1 | $17.8\pm 0.8$ | $14.30\pm 0.02$ | $-21.8\pm 0.5$
U | $21.6\pm 1.2$ | $14.32\pm 0.02$ | $-21.4\pm 0.4$
B | $36.3\pm 0.5$ | $15.57\pm 0.01$ | $-20.0\pm 0.3$
g | $38.3\pm 0.3$ | $15.18\pm 0.01$ | $-20.3\pm 0.3$
Table 3: The lightcurve decline rates in all optical bands at various phases of evolution. The epochs are given with respect to the estimated explosion epoch. | Time Interval
---|---
| 40-90d | 250-500d
| (mag d-1) | (mag d-1)
B | 0.009$\pm$0.002 | 0.011$\pm$0.001
g | 0.004$\pm$0.001 | 0.010$\pm$0.001
V | 0.002$\pm$0.001 | 0.012$\pm$0.000
r | -0.001$\pm$0.001 | 0.010$\pm$0.001
i | -0.004$\pm$0.001 | 0.009$\pm$0.001

The UV and optical lightcurves of ASASSN-14il are shown in Figure 3. The first
$\sim 40$ days after the explosion covers the initial rise in all UV and
optical bands. We perform a spline fit on the lightcurves to estimate the peak
magnitude and rise times. Monte-Carlo (MC) experiments were done to evaluate
the associated uncertainties with the fit. For each MC experiment, we draw
pseudo data points from the Gaussian distributions defined by the photometric
magnitude and the error bars of the original data points and perform the fit.
The rise time and peak magnitudes in UV-Optical bands are listed in Table 2.
We observe an increase in the rise time as we move towards the redder bands,
which is expected as the photosphere cools down and moves towards redder
wavelengths. The lightcurve in Swift UV bands reaches a maximum around day
15-18 and declines very sharply thereafter. The B and g band attains a maximum
at around day 36-38, much later than UV bands. The V and r bands do not have a
distinct maximum as the initial rising phase directly transitions into a very
flat plateau (day 40-100). The i band slowly rises in this duration. The
flattening of the optical lightcurves in optical bands between days 60-90
could be due to an ongoing CSM interaction. Both UV and optical lightcurves
are bumpy in nature, with the optical bands showing a prominent bump around
day 100. These small bumps in the lightcurve can be attributed to SN ejecta
encountering regions with enhanced CSM density (Graham et al., 2014). All
optical bands show a decline in late phases (day $>250$). The initial decline
of the lightcurves after the plateau is not well captured due to data gaps
(days 120-260). Therefore, it is not possible to comment on the similarities
with the transitional phase of SNe IIn-P. The lightcurve settles into the
decline of 1.0-1.2 mag (100d)-1 at late phases, which is close to the expected
56Co decay rate of 0.98 mag (100d)-1. It is possible to explain the observed
decay rate as a result of the ejecta-CSM interaction as well, as this
mechanism can give rise to a wide diversity of decay rates in SNe IIn. The
lightcurve decline rate in all optical bands at different phases of the
evolution is given in Table 3.

###  4.1 Comparison Sample

Table 4: General properties of the SNe IIn comparison sample. We have taken
redshift dependent distances assuming H0 = 73 km sec-1 Mpc-1,
$\Omega_{\mathrm{matter}}$=0.27, and $\Omega_{\mathrm{vaccum}}$=0.73. The
redshifts are corrected for the effects of GA, Virgo, and Shapley unless
mentioned otherwise. We take the reference epoch to be the explosion epoch if
available in the literature; otherwise, the discovery epoch is considered.

SN | Host Galaxy | Offset from center | Distance | Extinction | Discovery | Explosion | Reference†
---|---|---|---|---|---|---|---
|  |  | (Mpc) |  $E(B-V)$ (mag) | (mjd) | (mjd) |
SN 2005ip | NGC 290 | 2′′.8 E, 14′′.2 N | 33.7 | 0.047 | 53679.16 | - | 1
SN 2006gy | NGC 1260 | 0.′′941 W,0′′.363 N | 77.7 | 0.56 | 53996.3 | 53967.0 | 2, 3
SN 2006tf | SDSS J124615.80+112555.5 | 0′′.2 E, 0′′.7 N | 330 | 0.027 | 54081.0 | - | 4
SN 2009kn | MCG -3-21-6 | 17′′.75 E,15′′.27 N | 68.9 | 0.114 | 55130.5 | 55115.5 | 5
SN 2010jl | UGC 5189A | 2′′.4 E,7′′.7 N | 49.2 | 0.058 | 55503.3 | 55478.6# | 6, 7
SN 2011ht | UGC 5460 | 12′′.4 E,17′′.2 N | 20.5 | 0.061 | 55833.0 | - | 8
PTF11oxu | WISEA J033834.32+223242.7 | 0′′.06 E, 16′′.5 N | 376 | 0.176 | 55853.3 | 55837.3 | 9, NED
PTF11rfr | WISEA J014216.97+291625.6 | 0′′.19 W, 14′′S | 286 | 0.042 | 55906.1 | 55895.4 | 10, NED
SN 2012ab | SDSS J122247.61+053624.2 | 0′′.01 E, 0′′.70 N | 83.6 | 0.079 | 55957.9 | 55955.3 | 11
SN 2015da | NGC 5337 | 12′′.00 E, 14′′.00 N | 37.2 | 0.98 | 57031.9 | 57030.4 | 11, 12
ASASSN-15ua | GALEXASC J133454.49+105906.7 | 1′′.40E, 0′′.95S | 244.5* | 0.0578 | 57368 | - | 13
SN 2016aps | - | 0′′.15 | 1302* | 0.0263 | 57440.0 | - | 14
ASASSN-14il | WISEA J004532.52-141532.7 | 0′′.33 N, 0′′.26 W | 88.5 | 0.24 | 56931.1 | 56925.9 | This work

† References: (1) Stritzinger et al. (2012) (2) Smith & McCray (2007); (3)
Agnoletto (2009); (4) Smith et al. (2008); (5) Kankare et al. (2012); (6)
Jencson et al. (2016); (7) Zhang et al. (2012); (8) Mauerhan et al. (2013);
(9) Nyholm et al. (2020); (10) Gangopadhyay et al. (2020); (11) Tartaglia et
al. (2020); (12) Smith et al. (2023); (13) Dickinson et al. (2023); (14)
Nicholl et al. (2020)
$*$ Distances are not corrected for the effects of GA, Virgo and Shapley.
# The adopted explosion date is that of the first detection.

Figure 4: top: Absolute r/R-band lightcurve of ASASSN-14il compared with those
of the SNe IIn in the comparison sample. ASASSN-14il is very luminous only
dwarfed by SNe 2006gy and 2016aps. The lightcurve is very similar to SN
2015da, but the decline rate of ASASSN-14il is higher than SN 2015da and
similar to SNe 2006tf, 2016aps, and ASASSN-15ua. bottom: Intrinsic B-V color
of ASASSN-14il compared with the reference sample. The gradual reward
evolution of B-V color evolution is similar to SNe IIn.

We have chosen a fairly diverse comparison sample of eleven other well-
researched SNe. The sample consists of the following objects: SNe 1998S,
2005ip, 2006gy, 2006tf, 2009kn, 2010jl, 2011ht, PTF11rfr, PTF11oxu, 2012ab,
ASASSN-15ua, and 2016aps. These are strongly interacting SNe IIn, covering a
diverse range in luminosity, and show interaction signatures up to late times.
The basic parameters of the comparison SNe are laid out in Table 4. SNe 2005ip
and 2012ab represent the normal luminosity SNe IIn. SNe 2006gy, 2006tf,
2010jl, ASASSN-15ua, and 2016aps are all very luminous type SNe IIn having
peak M${}_{r/R}<-20$ that result from various CSM configurations. SNe IIn-P
2009kn and 2011ht, with a distinct plateau in their lightcurves, are also
included in this comparison sample. As noted by Nyholm et al. (2020), PTF11rfr
and PTF11oxu also show a plateau in their lightcurve, but unlike the SNe
IIn-P, they don’t show a sharp drop from the plateau into a 56Ni-dominated
tail. The lightcurve modeling of these SNe reveals a diversity in their CSM
configuration. SN 2006gy results from a massive compact CSM shell (Smith &
McCray, 2007), while more extended CSM structures power the comparatively
long-lived lightcurves of SNe 2006tf, 2010jl, and ASASSN-15ua (Smith et al.,
2008; Ofek et al., 2014; Dickinson et al., 2023).

We have compiled a sample containing the normal luminosity and very luminous
SNe IIn from the literature. Also, the different lightcurve shapes resulting
from different physical scenarios enable us to highlight the lightcurve
heterogeneity of SNe IIn and the role of CSM interaction in their photometric
evolution. The sample allows us to put the luminosity of ASASSN-14il in
perspective and infer the possible physical scenario by comparing and
contrasting their lightcurve and color evolution.

###  4.2 Absolute lightcurve and Color Evolution

Adopting the distance modulus $\mu=34.74\pm 0.15$ mag and total reddening
$E(B-V)=0.21\pm 0.08$ mag, we generate the absolute magnitude lightcurves of
ASASSN-14il. The $r$-band magnitude of ASASSN-14il plateaus at $\sim-20.2$ mag
over many tens of days. The $V$-band also shows similar behavior. Figure 4
shows the comparison of the absolute $r$-band lightcurve of ASASSN-14il with
other SNe from the comparison sample. ASASSN-14il stands out as one of the
most luminous SN in the comparison sample, which is in concordance with SNe
2006tf, ASASSN-15ua, 2010jl and is only dwarfed by the extremely luminous SNe
2006gy and 2016aps. The lightcurve shape of ASASSN-14il is very similar to SN
2015da except at the late epochs when SN 2015da declines much more slowly.
However, the decline rate of ASASSN-14il is similar to those of SNe 2006tf,
ASASSN-15ua, and 2016aps. The lightcurves of all these SNe were powered by
massive amounts (in order of tens of M⊙) of dense extended CSM structures that
allow the ejecta-CSM interaction to persist until late times (Smith et al.,
2008; Dickinson et al., 2023; Tartaglia et al., 2020; Smith et al., 2023;
Nicholl et al., 2020). The progenitors for these events that can provide the
required mass-loss are speculated to be massive LBVs. A similar progenitor and
physical scenario seems plausible for explaining the observed luminosities and
late-time decline rate of the lightcurves of ASASSN-14il as well.

SNe IIn-P 2009kn and 2011ht show a plateau in their lightcurve, but afterward,
they steeply drop into a radioactivity-dominated tail. On the other hand, SNe
like PTF11rfr and PTF11oxu show a plateau in their lightcurve, but their
lightcurves smoothly decline afterward in the $r$-band (Nyholm et al., 2020).
Due to the data gaps between days 120-260, it is not possible to establish
parallels with either of these classes of objects. However, we note that the
SNe IIn-P in our sample are generally much fainter than ASASSN-14il.

The B-V color evolution of ASASSN-14il is in agreement with the other SNe IIn
from the comparison sample. All the SNe show a redward evolution in the B-V
color as the photosphere cools down, ASASSN-14il maintains an overall bluer
color. The overall color evolution of ASASSN-14il is strikingly similar to SN
2015da.

After correcting the multi-band magnitudes obtained for extinction, the flux
integration (in UVW2, UVM2, UVW1, U, B, g, V, r, i bands) was performed with
blackbody corrections to attain the bolometric luminosities using SuperBol
(Nicholl, 2018). The lightcurve peaks at day $\sim$16 attaining a luminosity
of $1.6\times 10^{44}$ erg s-1. Gal-Yam (2012) suggested that SNe showing peak
magnitude less than –21 in any band or peak luminosity higher than $7\times
10^{43}$ erg s-1 may be considered as SLSNe. According to both these criteria,
ASASSN-14il can be considered as a SLSNe; however, this threshold is somewhat
arbitrary, as mentioned in Gal-Yam (2019), and it is unclear whether two
separate luminosity populations exist for SNe IIn.

Table 5: Fitting parameters for MOSFiT csm and csmni models. The shell and wind type CSM are represented by ‘s0’ and ‘s2’ suffixes, respectively. Model | Mej | vej | fNi | MNi | Mcsm | R0 | log $\rho_{0}$ | $\dot{M}$ | tcsm
---|---|---|---|---|---|---|---|---|---
| (M⊙) | (km s-1) | (%) | (M⊙) | (M⊙) | (AU) | (g cm-3) | (M⊙ yr-1) | (yr)
csm_s0 | $32.2^{+6.1}_{-5.7}$ | $4654.4^{+171}_{-162}$ | - | - | $9.1^{+1.9}_{-1.4}$ | $102.4^{+22.3}_{-17.2}$ | $-10.8^{+0.1}_{-0.1}$ | - | 4.8
csm_s2 | $108.4^{+5.8}_{-5.3}$ | $4685^{+97}_{-124}$ | - | - | $4.7^{+0.3}_{-0.3}$ | $19.8^{+4.2}_{-2.9}$ | $-10.8^{+0.1}_{-0.1}$ | 2.8 | 0.9
low_Ni+csm_s0 | $20.9^{+16.7}_{-12.8}$ | $4670^{+210}_{-200}$ | $0.4^{+0.3}_{-0.2}$ | $0.08^{+0.18}_{-0.07}$ | $15.87^{+12.53}_{-5.73}$ | $187.11^{+145.4}_{-68.2}$ | $-10.43^{+0.7}_{-0.3}$ | - | 8.9
low_Ni+csm_s2 | $14.6^{+2.7}_{-2.6}$ | $4897^{+172}_{-111}$ | $8^{+0.5}_{-0.5}$ | $1.2^{+0.3}_{-0.3}$ | $5.4^{+0.8}_{-1.0}$ | $80.2^{+13}_{-16}$ | $-10.36^{+0.11}_{-0.16}$ | 125.3 | 3.8
high_Ni+csm_s0 | $6.93^{+0.6}_{-0.5}$ | $5370^{+253}_{-241}$ | $66^{+5}_{-6}$ | $4.6^{+0.8}_{-0.7}$ | $1.9^{+0.7}_{-0.5}$ | $87.76^{+22.1}_{-14.5}$ | $-11.87^{+0.03}_{-0.02}$ | - | 4.2
high_Ni+csm_s2 | $7.6^{+0.6}_{-0.4}$ | $6025^{+284}_{-271}$ | $65^{+4}_{-5}$ | $4.9^{+0.7}_{-0.6}$ | $3.2^{+0.4}_{-0.3}$ | $126^{+16}_{-15}$ | $-12.1^{+0.1}_{-0.!}$ | 5.63 | 5.97

##  5 multi-band lightcurve modeling using MOSFiT

Figure 5: Results of csm models from MOSFiT. The lightcurves representative of
the evaluated posterior are plotted on top of the observational data. The
colored and gray lines represent the posterior of wind and shell csm models,
respectively. The lightcurves of ASASSN-14il are well reproduced by the CSM
models with Mcsm between 5-9 M⊙.

To estimate the physical parameters associated with the explosion, we also
performed analytical multi-band lightcurve modeling of ASASSN-14il. MOSFiT is
an open-source lightcurve fitter for transients (Guillochon et al., 2018). It
fits the multiband lightcurves with a list of built-in models using Markov
Chain Monte Carlo (MCMC) methods.

Since narrow and intermediate width lines, which are indicative of
interaction, are seen in the spectra of ASASSN-14il throughout its evolution,
we attempted to fit its lightcurves with the models which include CSM as a
powering mechanism. Chevalier (1982, 1998) lays the analytical groundwork for
the radiation expected from ejecta interacting with CSM. To achieve a self-
similar solution, it is assumed that the ejecta is expanding gas into a
stationary CSM. Both the expanding outer-ejecta and the stationary CSM are
assumed to follow a power-law distribution ($\rho_{\mathrm{ejecta}}\propto
r^{-n}$ and $\rho_{\mathrm{CSM}}\propto r^{-s}$). Chatzopoulos et al. (2012)
presented generalized semi-analytical models that take into account shock
power from interaction, magnetar spin-down, and 56Ni radioactivity. This work
has been used in Chatzopoulos et al. (2013) with $\chi^{2}$-minimization
algorithm to find optimal parameters for a sample of interacting SNe.
Chevalier (1982) presented the solutions for $s=0$ and $s=2$ and later Jiang
et al. (2020) extended these solution to $0\leq s\leq 2$. This model, along
with a generalized framework to easily fit multiband lightcurves, is
implemented in MOSFiT.

In the MOSFiT csm model, it is assumed that the SN luminosity results from the
conversion of kinetic energy from both the forward and reverse shock into
heating. While in the csmni model, 56Ni is considered as an additional power
source along with the csm interaction. We fitted the lightcurves of
ASASSN-14il with both the csm and csmni models. We fix the power law index of
the outer ejecta $n=12$, which is expected for the explosions of RSGs (Matzner
& McKee, 1999; Moriya et al., 2013). We fix the opacity $k=0.34$ cm2 g-1 for
fully ionized hydrogen-rich material. We take the rest of the parameters as
fitting parameters. The fitting was performed using the dynamic nested sampler
dynesty (Speagle, 2020). For ejecta mass and csm mass, we consider uniform
priors from 1-150 M⊙. For the inner radius of CSM (R0), we consider a uniform
prior from $1-500$ AU, and for the CSM density at the inner radius
($\rho_{0}$), we consider a log-uniform prior from $10^{-14}-10^{-9}$ g cm-3.
From both csm and csmni cases, we fitted with a shell csm ($s=0$) and a wind
csm ($s=2$) models. For the low_Ni and high_Ni models we consider a uniform
prior from 0.1-10% and 0.1-100% for the nickel fraction in ejecta (fNi = MNi /
Mej), respectively. The estimated parameters from fitting various models are
listed in Table 5. The $1\sigma$ errors listed represent the statistical
uncertainty associated with fitting and hence are not indicative of the model
uncertainties.

Both the csm and csmni models (with either a wind or shell CSM configuration)
provide a very similar fit to the observed lightcurves, despite the latter
being the more complex. Both are able to reproduce the broad lightcurve
features reasonably well. However, these models have a smooth CSM
distribution, and therefore they are unable to reproduce both the lightcurve
bumps and the decline of lightcurve in the UV bands after the peak. We tried
to model the lightcurve after removing the bump seen around day 90, in a
manner similar to Hosseinzadeh et al. (2022), but this approach doesn’t
improve the fit to the lightcurves in our case. Figure 5 shows the sample
lightcurves for the csm (shell and wind) models, which are representative of
the posterior for these models. The late-time luminosity is powered by the
reverse shock running through the ejecta after the forward shock terminates.
The models having an extremely high fraction of 56Ni in the ejecta can
reproduce the steep decline of UV lightcurves. But due to their unphysical
nature, we neglect these models. The variation of the estimated parameters for
different models dominates over the fitting uncertainties. However, the
requirement of a high ejecta mass (10–100 $M_{\odot}$) and a very dense CSM
($\sim 10^{-11}$ g cm-3) are consistent.

The mass-loss rate for the wind-like csm models is calculated as $\dot{M}=4\pi
v_{w}\rho_{0}R_{0}^{2}$, and the time of ejection of CSM (tcsm) is calculated
as $R_{0}/v_{w}$, where we have assumed a typical LBV-like wind velocity
($v_{w}$) of 100 km sec-1. We estimate a mass-loss rate of 2.8 M⊙ yr-1 (for
the csm_s2 model), and the CSM shell/wind expelled about 1-5 years prior to
the explosion.

##  6 Spectroscopic Evolution

Figure 6: Spectral Evolution of ASASSN-14il from day 8–327. Spectra are
displayed in the source rest frame and are corrected for extinction. The grey
color represents the original spectra, and the colored lines represent
smoothed spectra. Early spectra have a blue continuum and narrow H and He
lines with Lorentzian wings. Broad ejecta components can be seen in H$\alpha$
and H$\beta$ lines along with intermediate width lines at later epochs. The
host spectrum is plotted at the bottom.

Figure 7: Comparison of ASASSN-14il spectra with the comparison sample at
distinct stages of its evolution. SNe IIn-P and SNe IIn are plotted above and
below the spectra of ASASSN-14il (black). ASASSN-14il shows prominent Balmer
features from early times similar to other SNe IIn and SNe IIn-P. At mid
epochs, the Balmer lines are much more prominent in the SNe IIn-P owing to
their weaker continuum, in contrast to the SLSNe. At late epochs, many of the
SNe develop a noticeable blueshift in their H$\alpha$ profiles. Figure 8: NIR
spectral evolution of ASASSN-14il. The NIR spectra are plotted along with
optical spectra from nearby epochs. The spectra are displayed in the source
rest frame and are not corrected for reddening. The gray color represents the
original spectra, and the colored lines represent smoothed spectra. The
evolution of Pashcen lines is similar to the optical Balmer lines.

The spectral evolution of ASASSN-14il from day $\sim$8–327 is shown in Figure
6. From the first spectra, narrow line features of Hydrogen are visible, which
contrasts with some SNe IIn, where initial spectra are almost featureless. The
overall spectral evolution is very slow. Multicomponent H$\alpha$ line is
visible throughout its spectral evolution. Other Balmer lines, as well as weak
He lines (5876, 6678, and 7065 Å) are also present since the beginning. The
evolution of H$\alpha$ is very smooth, and a detailed description is given in
subsection 6.2. The lines originating from SN ejecta like Ba II, Sc II, Mg I,
Ti II, and Ca II that are usually seen in a normal CCSNe without predominant
interaction are not visible in the early phase of spectral evolution. Some
lines like O III and S II are seen in the spectral evolution, indicating
contamination by the host galaxy. For comparison, we have also shown the host-
galaxy spectrum in Figure 6. The host galaxy spectrum was taken on UT
2004-09-06.61 under the 6dF galaxy survey (Jones et al., 2004, 2009).

The first spectrum was taken on day 8, featuring a blackbody continuum at a
temperature of 11,000 K and a radius of 100 AU. This spectrum shows narrow
H$\alpha$ profiles on top of smooth Lorentzian wings. Similar Lorentzian wings
are present in other Balmer and He lines. The N II contamination from the host
galaxy is also seen in this spectrum.

The spectra are dominated by Lorentzian line profiles and a smooth blue
continuum up to day 30. At day 50, the line profile slowly shifts from
Lorentzian into more complex multi-component profiles. During this time, the
temperature of the continuum decreases to 6000 K and becomes constant
thereafter. After day 50, ejecta signatures are visible as a broad component
in the H$\alpha$ line.

After day 100, a clear blueshift can be noticed in the H$\alpha$ profile, the
extent of which increases with time. At day 327, the Balmer lines are
dominated by the broad ejecta component, and CaII (both forbidden and NIR)
lines are also visible. The blueshift seen earlier is even more prominent at
this stage. The nebular phase spectral lines typical for the CCSNe are not
visible at these later stages either. The systematic blueshift in the
H$\alpha$ profile can arise due to the asymmetric CSM or dust formation in the
post-shock gas (Smith et al., 2008; Jencson et al., 2016).

The narrow emission feature of the HeI 5876 Å line is noticeably redshifted
compared to the rest-wavelength in the higher resolution ANU spectra (taken
between day 8-32; Figure 2). Such a redshift can be caused by an unresolved
P-cygni profile characteristic of a CSM outflow (Smith et al., 2023). The
blue-edge of the narrow P-Cygni in SN 2017hcc lies at -70 to -60 km s-1, which
is completely washed out in low to moderate resolution spectra (Smith &
Andrews, 2020). ASASSN-14il maybe a similar case where the P-Cygni profile in
HeI and corresponding narrow P-cygni profiles in Balmer lines are unresolved
due to the limited resolution of the spectra. It is also possible that these
narrow P-Cygni features are masked due to contamination from the host.

Comparison of ASASSN-14il spectra with other SNe IIn from the comparison
sample at different stages of its evolution is shown in Figure 7. All the SNe
have a blue continuum in the early spectrum (day $\sim 20$). All interacting
SNe show narrow emission lines superimposed on the continuum. The H$\alpha$ of
ASASSN-14il is very prominent, similar to most of the SNe in the comparison
sample signifying strong ongoing interaction. However, SN 2012ab exhibits
relatively weaker H$\alpha$, which results from directly seeing the high-
velocity jet-like ejecta with relatively lesser CSM interaction at this phase
(Gangopadhyay et al., 2020). The early spectra of ASASSN-14il bear an overall
resemblance to the other events in the sample except for SN 2005ip, which
shows a spectrum similar to normal type II, with narrow lines from preshock
CSM on top (Smith et al., 2009b).

The spectral evolution of ASASSN-14il at mid-phases closely resembles SNe
2006gy and 2015da around day $\sim$ 120\. SN 2010jl has similar H$\alpha$ line
morphology but with more prominent Balmer lines than ASASSN-14il. This
spectral synergy between the three SNe is also evident in their photometric
behavior, showing high luminosities and long-lasting lightcurves. In contrast,
for the SNe IIn-P 2009kn and 2011ht, H$\alpha$ line emission is more prominent
with respect to their weaker continuum. In addition to the narrow component, a
broad component can be clearly seen in ASASSN-14il and all the other SNe IIn.
This component arises due to the ejecta contribution in the H$\alpha$ profile.
Additionally, the broad component of ASASSN-14il shows a deficit in the red
side flux of the H$\alpha$ line similar to SNe 2005ip, 2010jl, and 2015da.

After $\sim 1$ year of the explosion, ASASSN-14il shows a flat-topped broad
component in H$\alpha$, with the narrow component on top. The observed
blueshift in ASASSN-14il and SNe 2005ip, 2010jl, and 2015da is much more
prominent compared to the mid-phases. The overall blueshift in the line
profile of SN 2010jl was interpreted as the broadening due to electron
scattering (Fransson et al., 2014). However, Smith et al. (2012, 2023) argue
that broadening by electron scattering should be symmetric about the original
source of the narrow-line photons. The narrow lines in SNe 2010jl and 2015da
are found at rest-frame velocities even at late phases where a blueshift is
seen in the H$\alpha$ profile. Instead, this blueshift is explained as a
result of dust formation in the post-shock CSM/ejecta, which preferentially
blocks the emission from receding material, which is consistent with the other
evidence of dust-formation (Maeda et al., 2013; Gall et al., 2014). The
H$\alpha$ profile of ASASSN-14il also shows an overall blueshift and a rest-
frame emission component. The H$\alpha$ profile of ASASSN-14il at this epoch
is similar to SN 2005ip at day $\sim 177$, which shows a similar blueshifted
flat top profile. This blueshifted profile appears in SN 2005ip as early as
day $\sim 120$ and is attributed to the post-shock dust formation in the CDS
(Fox et al., 2009; Bak Nielsen et al., 2018). Based on these comparisons, the
blueshifted H$\alpha$ profile of ASASSN-14il is likely the result of dust
formation in the post-shock CSM/ejecta.

###  6.1 Infrared Spectrum

Five NIR spectra were obtained for ASASSN-14il between days 19–116 after the
explosion. They are displayed in Figure 8 along with the closest optical
spectra. Prominent lines of Pa$\alpha$, Pa$\beta$, and Pa$\delta$ can be seen
in the spectral sequence of ASASSN-14il. Pa$\gamma$ is also visible but is
usually blended with He i 10830 Å and Pa$\epsilon$ line. The FWHM of the
Pa$\beta$ and Pa$\delta$ typically varies between 8,000 - 13,000 km sec-1. The
lines here are isolated, and FWHMs concord with H$\alpha$ within error bars.
Overall, the NIR spectra of ASASSN-14il are consistent with optical spectra
regarding interaction features, mostly implied from the FWHM’s being
consistent in both the wavelength regimes.

###  6.2 H-alpha Evolution

ASASSN-14il suffers significantly from host galaxy lines of OII, NII, and SII
(seen clearly in the higher resolution WiFeS spectral sequence). The host-
galaxy spectrum taken under the 6dF survey (see section 6) has an FWHM
resolution of 9–12 Å in the H$\alpha$ region. The SN spectra have lower
resolution compared to the host spectrum. The SN spectrum was subtracted after
degrading the resolution of the host galaxy spectrum, taking the SII 6717-6731
Å line as a reference for host contribution. In most low to medium-resolution
spectra, the host galaxy contribution is almost removed, leaving an unresolved
narrow residual component behind. For the ANU-WIFES spectra, the analysis was
performed as is, without performing the host-galaxy subtraction, as they are
much higher in resolution than the available host spectra.

Figure 9: The H$\alpha$ evolution of ASASSN-14il being fitted with multiple Lorentzian and Gaussian components. The host galaxy contribution is subtracted, and the continuum is selected by masking the line regions. ASASSN-14il shows a very complex geometry of the spectral profiles. The orange, blue, and green components represent the narrow, intermediate, and broad components, respectively. For day 31.6 and 327.4 The H-beta (violet) from is plotted in violet. Table 6: This table describes the spectral decomposition of H$\alpha$ profile of ASASSN-14il. The line centers and FWHMs of the narrow, intermediate-width, and broad components are reported. Phase | Source | Narrow Component | Intermediate-width Component | Broad Component
---|---|---|---|---
|  | center | fwhm | center | fwhm | center | fwhm
(d) |  | (km s-1) | (km s-1) | (km s-1) | (km s-1) | (km s-1) | (km s-1)
7.7 | ANU | 131.5$\pm$0.3 | 190.6$\pm$0.9 | 96.7$\pm$9.6 | 2039.4$\pm$27.6 | - | -
7.8 | LCO | -128.1$\pm$5.3 | 542.2$\pm$18.2 | -54.0$\pm$10.2 | 1890.1$\pm$55.6 | - | -
9.7 | LCO | -23.0$\pm$4.7 | 677.3$\pm$16.6 | 18.2$\pm$10.0 | 2196.4$\pm$61.3 | - | -
11.8 | LCO | 37.7$\pm$8.4 | 628.7$\pm$28.6 | 113.4$\pm$24.5 | 2360.8$\pm$129.1 | - | -
13.6 | ANU | 148.3$\pm$0.3 | 190.8$\pm$0.8 | 144.9$\pm$8.0 | 2112.6$\pm$22.7 | - | -
14.6 | LCO | 28.7$\pm$6.1 | 701.6$\pm$21.8 | 58.2$\pm$14.3 | 2239.2$\pm$89.5 | - | -
17.7 | LCO | 72.5$\pm$17.3 | 568.0$\pm$58.6 | -30.9$\pm$26.6 | 1680.0$\pm$151.3 | - | -
22.7 | ANU | 136.7$\pm$0.3 | 198.9$\pm$0.9 | 134.4$\pm$8.3 | 1949.5$\pm$24.1 | - | -
22.8 | LCO | -9.4$\pm$4.1 | 755.4$\pm$12.1 | 153.8$\pm$27.3 | 3837.7$\pm$79.5 | - | -
26.1 | PESSTO | 193.6$\pm$8.5 | 460.0$\pm$28.7 | 89.1$\pm$11.0 | 1784.7$\pm$51.9 | - | -
26.6 | LCO | 38.1$\pm$5.0 | 688.4$\pm$17.7 | 92.2$\pm$12.9 | 2090.0$\pm$82.1 | - | -
31.6 | ANU | 144.7$\pm$0.3 | 178.1$\pm$0.9 | 130.2$\pm$8.2 | 1701.7$\pm$24.7 | - | -
49.2 | PESSTO | 239.6$\pm$4.8 | 497.3$\pm$17.3 | 118.3$\pm$12.8 | 1549.7$\pm$57.1 | 339.0$\pm$69.6 | 5010.3$\pm$243.0
81.6 | LCO | 206.8$\pm$1.6 | 506.3$\pm$5.6 | 174.3$\pm$6.3 | 1841.0$\pm$28.9 | -973.3$\pm$70.5 | 8680.9$\pm$241.4
109.2 | PESSTO | 66.3$\pm$7.6 | 715.4$\pm$29.4 | -59.5$\pm$29.3 | 2098.5$\pm$130.8 | -573.2$\pm$57.3 | 8197.7$\pm$192.6
123.1 | PESSTO | - | - | 44.0$\pm$27.9 | 1531.5$\pm$82.9 | -461.2$\pm$52.4 | 7608.1$\pm$155.3
327.4 | PESSTO | -39.6$\pm$9.7 | 526.5$\pm$24.5 | - | - | -957.0$\pm$68.7 | 6543.2$\pm$165.0

In the host-galaxy subtracted spectra, the H$\alpha$ profile was fitted with
combinations of the following to reproduce the overall line profile - i) A
narrow Gaussian component, ii) A Lorentzian/Gaussian intermediate width
component, iii) A broad Gaussian component. At early times, the overall
H$\alpha$ profile is better represented by a Lorentzian intermediate width
component as the emission lines are dominated by electron scattering. At late
times (day $>$50), a Gaussian intermediate width component better reproduces
the overall line profile as the lines evolve into a more complex multi-
component structure. The choice of the continuum is very critical for
performing these fits. The continuum is selected to be far from the line
region by at least 50 Å. The continuum selection was performed repetitively to
optimize the consistency of the fits. The spectral evolution and the
corresponding fits at selected epochs are shown in Figure 9. The parameters of
components obtained from the fitting at all epochs are given in Table 6. The
errors listed represent only fitting errors, and other uncertainties like
resolution matching, subtraction of host-spectra, and imperfections in
wavelength calibration have not been included. The evolution of the line
centers and FWHMs of the various components is plotted in the top panel of
Figure 10.

Three components can describe the line profiles: i) A narrow component
(unresolved in the lower resolution spectra) centered around the rest
wavelength, ii) An intermediate width component centered around the rest
wavelength, with FWHM typically varying between 1600–2000 km sec-1, and iii) A
broad emission varying in FWHM between 5000–8500 km sec-1 seen after day 50.
This broad component progressively develops a blueshift of $\sim 600-1000$ km
s-1. The minimum FWHM of the narrow component is 178 km sec-1 seen in the ANU
spectra (day 31.6), however it still may be unresolved. To check whether the
narrow component seen in the high-resolution ANU spectrum is due to the host
or intrinsic to ASASSN-14il, the host-galaxy spectrum was subtracted by
degrading the resolution of the ANU spectrum. A residual narrow component
persists, indicating narrow line emission from the SN as seen in our spectral
sequence. Dickinson (2021) performed a similar analysis with the high-
resolution spectrum taken after the SN had faded and arrived at a similar
conclusion. It should be noted that this residual narrow component may still
be from the host galaxy as the H to N emission ratio is variable depending on
the slit positions and orientation. However, a strong narrow component is seen
in the HeI $5876$ Å line as well, in spite of the host galaxy having very low
HeI emission.

Figure 10: Top Panel: The FWHMs of the different components of the H$\alpha$
obtained by fitting Gaussian/Lorentzian to line profiles. The three color-
coded zones in the panels demarcate the regions where ASASSN-14il enters into
three phases of its evolution as described in subsection 6.2; Middle Panel:
the line flux ratios of H$\alpha$ and H$\beta$ for long-lasting SNe IIn.
10$\%$ error has been included in the flux-ratios; Bottom Panel: the line-
luminosity variations of H$\alpha$ as compared with other members of the sub-
class.

The H$\alpha$ evolution of ASASSN-14il can be divided into three distinct
phases. From day 8–50, the H$\alpha$ is dominated by Lorentzian wings caused
by the electron scattering of narrow line photons. A narrow component is also
visible in the spectra even after subtracting the host galaxy contribution.
This indicates that the photosphere lies in the unshocked CSM at this phase.
The UV (and bolometric) lightcurves peak during this phase (around day 15),
which can ionize the unshocked CSM. During this phase, no signatures of
underlying ejecta are seen in the H$\alpha$ profile, unlike some SNe where
broad emission or broad P-Cygni profiles are seen from early times, which
would indicate direct line-of-sights to the ejecta. In SN 2012ab, this is
achieved by disk-like CSM (Gangopadhyay et al., 2020) and in SN 2005ip, a
clumpy CSM structure is present (Smith et al., 2009b).

From day 50–80, the combination of a narrow Gaussian and a Lorentzian profile
can no longer accurately reproduce the line profiles. Instead, the line
profiles are more faithfully represented by combining three Gaussian (narrow,
intermediate width, and broad) components. The broad wings of high-velocity
ejecta are distinguishable in the line profiles. The Gaussian intermediate
width and broad with components have FWHMs of 1550–2000 km s-1 and 5000–8000
km s-1, respectively. No significant P-Cygni absorption associated with the
broad H$\alpha$ component is visible, suggesting that the underlying ejecta is
already entering the optically thin regime.

After day 81.5, both the intermediate width and broad component are visible up
to day 123.0, along with the narrow component present throughout the
evolution. During the evolution, the broad component changes in FWHM between
5000—8800 km s-1. However, after day 109, a prominent blueshift can be noticed
in the broad component of H$\alpha$, now centered between $\sim$ -500 to -900
km s-1. The H$\alpha$ profile in the day 327.3 spectrum can not be described
as a combination of Gaussians with shifted centroids. Instead, it can be
better described as missing flux at velocities higher than -4000 km s-1.
Figure 9 also shows the H$\beta$ line plotted (with offset for clarity)
together with H$\alpha$ at days 31.6 and 327.4. The H$\beta$ line is scaled
such that it matches the blue-side profile of the H$\alpha$ line. With this
comparison, it is clear that the blueshift seen at late epochs is wavelength-
dependent, more prominent at shorter wavelengths. This is further discussed in
section 7.

Figure 10 shows the flux ratio of H$\alpha$ to H$\beta$, line luminosities of
H$\alpha$ as compared with some other superluminous and long-lasting
interacting SNe IIn (SN 2005ip – Smith et al., 2009b; SN 2006tf – Smith et
al., 2008; SN 2010jl –Jencson et al., 2016; SN 2015da –Tartaglia et al., 2020
and ASASSN-15ua – Dickinson et al., 2023) along with FWHMs of the H$\alpha$
components of ASASSN-14il. The H$\alpha$ luminosity includes the contribution
from the narrow, intermediate-width, and broad components of the H$\alpha$
profile. The H$\alpha$ line fluxes and H$\alpha$ to H$\beta$ flux ratios were
estimated from the dereddened spectra, including a 10$\%$ error. As seen in
the middle panel of Figure 10, the H$\alpha$ to H$\beta$ flux ratio of
ASASSN-14il steadily increases throughout its evolution. Dickinson et al.
(2023) have proposed that the H$\alpha$ to H$\beta$ flux ratio at early epochs
is dominated by recombination after photoionization, where the line ratio
typically approaches 3. Then, at later times, it is dominated by heating by
the post-shock gas dominated by collisional excitation. A similar trend is
noticed in ASASSN-14il and other objects where a transition is seen from
photoionization to collisional excitation. The flux ratio of ASASSN-14il also
closely resembles SN 2015da throughout the evolution.

The bottom panel of Figure 10 compares the H$\alpha$ line luminosities of
ASASSN-14il with other SNe IIn. For ASASSN-14il, the luminosity initially
increases from day 8–23 by a factor of 7; the luminosity is almost constant
from day 23–82, and from day 82–327, it increases again by a factor of 3. The
H$\alpha$ luminosity of SNe 2010jl, 2015da and ASASSN-15ua shows an increasing
trend when the interaction becomes maximum and decreases later. ASASSN-14il
shows a similar trend, except for the flattening seen between days 23–82,
which could be due to persistent interaction with shell/clump.

Figure 11: Cartoon diagram representing the evolution of ASASSN-14il through
various distinct phases. Representative line-of-sights (LOSs) are shown by
solid lines, where dashed lines indicate hindered LOSs. A dense disk-like CSM
is located face-on to the observer, and A low-density CSM is present
elsewhere. Interaction with dense CSM gives rise to most of the luminosity. At
mid-phases, the photosphere recedes from the pre-shock CSM and signatures from
the post-shock region and ejecta are visible. At late times the formation of
dust obscures the tradition from receding material.

##  7 Results and Discussion

This paper presents the long-term photometric (up to day 480) and
spectroscopic (up to day 327) monitoring campaign of a SN IIn ASASSN-14il. The
lightcurves show bumpy behavior at different stages, which can be attributed
to ejecta interacting with regions of enhanced CSM density. The multi-band
lightcurves also display a plateau-like behavior due to ongoing interaction. A
similar plateau is seen in SNe IIn-P (Mauerhan et al., 2013) followed by a
steep drop (2-4 mag in optical) from the plateau into a radioactivity-
dominated tail. In contrast, some SNe IIn show a plateau followed by a linear
decline (Nyholm et al., 2020). However, we can not ascertain the nature of
ASASSN-14il in the context of these two classes due to gaps in the
observations. ASASSN-14il, with a peak absolute magnitude Mr $\sim-20.3\pm
0.2$ mag, is one of the brightest SN IIn. In UV bands, it is brighter than
$-21$ mag, which satisfies the criteria for SLSN category (Gal-Yam, 2012).
However, whether such bright interacting transients form a separate population
or fall in the tail end of normal SNe IIn is unclear. The color curves are
concordant with a typical SN IIn. Overall, it bears a stark resemblance with
SN 2015da in terms of the lightcurve and color evolution which indicates
similar powering mechanism(s).

Both wind and shell type CSM models in MOSFiT can reproduce the broad
lightcurve features of ASASSN-14il. To reproduce the observed luminosities of
ASASSN-14il, a high amount of ejecta mass (a few tens to hundred M⊙) is
required to interact with a high-density CSM ($\sim$ 10-11 g cm-3). These one-
zone models, however, can not reproduce the steep decline of UV lightcurves
and the small bump seen in the optical lightcurves around day 90. The
assumptions of homologous expansion, constant opacity, and spherical symmetry
are simply not representative of the actual explosion scenario. More
sophisticated modeling techniques are required to adequately model the
lightcurve, which is beyond the scope of this work.

The spectral evolution of ASASSN-14il resembles typical SNe IIn dominated by
smooth blue continuum and Balmer lines. The Balmer lines are dominated by
Lorentzian wings in the early epochs. The H$\alpha$ line profile evolves to be
more complex and multicomponent after day 50. The signs of underlying high-
velocity ejecta are visible after this epoch. The emergence of ejecta
signatures at the time of peak optical luminosity indicates asymmetries in the
explosion/CSM geometry, which is quite common in SNe IIn (Smith, 2017;
Bilinski et al., 2024). This is further discussed in subsection 7.2.

After day 123.0, a prominent overall blueshift can be observed in the
H$\alpha$ profile, which becomes more pronounced in the final spectra at day
327. The spectra of ASASSN-14il show similarity with SNe 2006gy, 2010jl, and
2015da. The late-time flat-topped profile is reminiscent of SN 2005ip. Also,
the late-time blueshift can be explained by dust formation in the post-shock
CSM or ejecta (similar to SNe 2005ip, 2010jl, and 2015da; Smith et al., 2009b,
2012, 2023). The case of dust formation and alternative explanation are
discussed in subsection 7.4. The evolution of H$\alpha$ to H$\beta$ flux ratio
supports the transition scenario from recombination to collisional excitation.
The line luminosity of H$\alpha$ shows an increasing trend due to persistent
ongoing interaction.

###  7.1 Narrow lines

ASASSN-14il shows narrow lines throughout its spectral evolution. However, the
narrow lines do not show a P-Cygni profile which is expected from a CSM
outflow. The lack of a narrow-width P-Cygni feature is very uncommon among SNe
IIn. In the case of ASASSN-14il, there could be several reasons for the lack
of this feature. First, it could be just due to the spectral resolution of the
available spectra. P-Cygni profiles corresponding to outflow velocities of
$<100$ km s-1 would be unresolved in the presented spectra. An unresolved
P-Cygni profile will appear redshifted as explained by Smith et al. (2023),
and there are indeed some hints of it as the narrow HeI 5876 feature is
redshifted from its expected rest wavelength. It is possible that ASASSN-14il
is surrounded by an asymmetric CSM structure such that the CSM along the line-
of-sight has a relatively lower density (as well as slow-moving) which will
give rise to a weaker P-Cygni feature that is hard to resolve under limiting
conditions. Secondly, it could be due to the contamination of the narrow lines
by the host-galaxy emission.

Additionally, a recent study by Ishii et al. (2024) explores the relations
between the line shapes and CSM structure by Monte Carlo radiative transfer
codes. They find that a narrow line exhibits a P-Cygni profile only when an
eruptive mass-loss event forms the CSM. The CSM structure from a steady mass
loss will have a negative velocity gradient after the SN event due to
radiative acceleration. Therefore, an H$\alpha$ photon emitted at the deeper
CSM layers, traveling outwards, will never be able to undergo another
H$\alpha$ transition. Therefore, the CSM residing in the direction of the
observer may be created by steady-mass loss (e.g., LBV winds) in the case of
ASASSN-14il.

###  7.2 Asymmetry

In a spherical symmetric model for interacting SNe at early times the
continuum photosphere lies in the CSM ahead of the shock. After some time, the
photosphere moves into the post-shock CDS, and the uninterrupted ejecta is
visible only at late times when the interaction has slowed down enough (Smith,
2017). The line profile of H$\alpha$ earlier spectra (up to day 31) of
ASASSN-14il show a symmetric Lorentzian profile expected from the continuum
being in the pre-shock CSM. However, the broad component from the SN ejecta is
visible as soon as day $50$, along with the intermediate component from the
post-shock CDS. The optical luminosity is near-peak at this epoch, indicating
the presence of strong CSM interaction. This is not explained from a
spherically symmetric CSM geometry. Instead, the emergence of broad ejecta
suggests significantly lower CSM along the LOS compared to the CSM needed to
sustain the peak luminosity. This can be explained by an asymmetric CSM
configuration such that dense CSM, at angles alternate to the LOS, giving rise
to the majority of the observed luminosity. Such a configuration is not
uncommon in SNe IIn. SNe 2010jl, 2012ab, and 2015da all reveal the underlying
ejecta earlier than expected from a spherical CSM distribution and are
suggested to feature a disk-like/torus-like CSM geometry (Katsuda et al.,
2016; Bilinski et al., 2020; Gangopadhyay et al., 2020; Smith et al., 2023).

A recent spectropolarimetry study of SNe IIn (including ASASSN-14il) by
Bilinski et al. (2024) reveals a significant amount of polarization in the
majority of events in their sample, indicating asymmetries in the CSM. They
study ASASSN-14il at 3 epochs (day 31, 63, 119 from the explosion date
estimated in this work). There is a change in continuum polarization going
from day 31 to 63, which coincides with the transition of the continuum
photosphere from the pre-shock CSM to CDS/ejecta. The continuum polarization
of ASASSN-14il is one of the lowest (along with SN 2014ab) in the sample.
Additionally, there is no significant line depolarization noticed in
ASASSN-14il. This polarization doesn’t necessarily eliminate the possibility
of asymmetric CSM; instead, they can be explained in a complimentary manner. A
disk-like CSM viewed from the polar region will have a face-on symmetry for
the continuum photosphere and, therefore, low-polarization (Bilinski et al.,
2020).

###  7.3 Mass-loss Rate

The intermediate-width component seen throughout the evolution of ASASSN-14il
indicates the persistent ejecta-CSM interaction, as also observed in the case
of SN 2012ab (Gangopadhyay et al., 2020). Assuming that the luminosity of the
ejecta-CSM interaction is fed by energy at the shock front, the progenitor
mass loss rate $\dot{M}$ can be calculated using the relation of Chugai &
Danziger (1994):

$\dot{M}=\frac{2L}{\epsilon}\frac{v_{w}}{v_{SN}^{3}}$ (1)

where $\epsilon$ ($<1$) is the efficiency of conversion of the shock’s kinetic
energy into optical radiation (an uncertain quantity), $v_{w}$ is the velocity
of the pre-explosion stellar wind, $v_{\mathrm{SN}}$ is the velocity of the
post-shock shell, and $L$ is the bolometric luminosity of ASASSN-14il. We
assume typical unshocked wind velocity observed for LBV winds $v_{w}$$\sim$100
km sec-1. The shock velocity is inferred from the intermediate width component
after the emission lines are no longer dominated by electron scattering (day
$>=$ 50). We take the shock velocity to be 1750 km s-1, similar to the value
used by Dickinson (2021). Using the bolometric luminosity at day $\sim 50$
(L=4.7 x 1043 erg sec-1) and assuming 50$\%$ conversion efficiency with
$\epsilon$=0.5, the estimated mass-loss rate for ASASSN-14il is 5.6 M⊙ yr-1.
However, we find that the estimated mass-loss rate is 1.0 M⊙ yr-1 for only the
integrated observed luminosity with no bolometric corrections. Dickinson
(2021) found the mass-loss rate to be 1 M⊙ yr-1 following similar analysis
based on the V-band luminosity. These values are also comparable to the mass-
loss value for the csm_s2 model derived from lightcurve fitting.

These values are comparable, albeit slightly higher, than the estimates for SN
2015da (Smith et al., 2023) at similar epochs. This can be attributed to
higher bolometric luminosity and slower shock speed in ASASSN-14il. The
estimated value of mass-loss rate is much higher than the typical LBV winds
(Smith, 2014) and higher than the values often attributed to some SNe IIn,
which are of the order of 0.1 M⊙ yr-1 as observed in some giant eruptions of
LBVs (Chugai et al., 2004; Kiewe et al., 2012). These values are also much
higher than those of normal-luminosity SNe IIn like SN 2005ip ($2-4\times
10^{-4}$ M⊙ yr-1; Smith et al., 2009b). It is also much larger than the
typical values of RSG and yellow hypergiants ($10^{-4}-10^{-3}$ M⊙ yr-1;
Smith, 2014), and quiescent winds of LBV (10${}^{-5}-10^{-4}$ M⊙ yr-1, Vink,
2018). The obtained mass-loss rate in ASASSN-14il indicates the probable
progenitor to be an LBV star that underwent an eruptive phase. The CSM may be
the result of interaction with a binary companion, which in would explain the
expected asymmetry in the geometry.

###  7.4 A case for Dust Formation

The H$\alpha$ profile of ASASSN-14il shows a deficit in the flux of the
redside wing. This deficit is clearly visible as an overall blueshift of the
line after day 100 and increases with time. At day 327, even a combination of
blueshifted Gaussian components can not appropriately reproduce the H$\alpha$
line profile. At this epoch, the line profile can be best described as missing
flux from the redside wing and near rest velocity. The flux deficit phenomenon
is wavelength dependent as well, affecting the shorter wavelength H$\beta$
more severely compared to H$\alpha$.

An overall blueshift could be caused by many reasons, such as radiative
acceleration, lopsided SN explosion or asymmetric CSM, obscuration of the
receding material by the continuum photosphere, and dust formation. However,
except for dust formation, none of the other mechanisms can explain the
observed time dependence and wavelength dependence.

In the case of a blueshift caused by the radiative acceleration of the CSM,
the blueshift should decrease as the luminosity drops, and we expect no
significant wavelength dependence for this case. Additionally, in the case of
a blueshift caused by the radiative acceleration scenario, the original narrow
line photon source should be blueshifted as well, which is not the case for
ASASSN-14il (Dessart et al., 2015). Similarly, in the case of obscuration by
the continuum photosphere, the expected blueshift will be strongest in early
times and decrease later on as the continuum optical depth drops. For a
lopsided explosion/CSM any blueshift present should be present from early
times as well, remaining consistent throughout the evolution. Thus none of
these scenarios are consistent with the observed evolution of ASASSN-14il.

However, all these requirements are readily explained by a scenario where new
dust grains form in the ejecta or the post-shock CSM. The dust formation will
increase at late times as the ejecta and the post-shock CSM cool down,
explaining the time evolution of the observed deficit of flux in the side wing
of H$\alpha$. Also, the wavelength dependence of the flux deficit is a natural
consequence of the extinction from dust affecting shorter wavelengths more
prominently. Dust formation has other signatures, such as an increase in the
NIR flux and an increased rate of fading the optical flux. However, due to
data gaps between days 120–250 and the lack of late-time NIR data, we are
unable to comment on it. Dust formation is very common in SLSNe IIn as seen in
the case of SNe 2006tf (Smith et al., 2008), 2010jl Smith et al. (2012); Maeda
et al. (2013); Gall et al. (2014), 2015da Smith et al. (2023), ASASSN-15ua
Dickinson et al. (2023), and 2017hcc Smith & Andrews (2020), so expecting dust
formation in ASASSN-14il would not be unreasonable.

###  7.5 Physical Scenario

Figure 11 describes a physical scenario that, although not unique, explains
the observables of ASASSN-14il through various stages of its evolution. In
this scenario, the ejecta from the SN explosion interacts with a disk-like
CSM. The viewing angle of the observer is such that the CSM disk is face-on to
the observer. This configuration is surrounded by distant CSM that may result
from steady mass loss from the progenitor.

Figure 11a describes a scenario where the photosphere lies in the pre-shock
CSM surrounding the interaction region. The line profiles are dominated by
narrow lines and Lorentzian wings resulting from electron scattering in the
ionized CSM. The signatures from the SN ejecta are masked by the interaction
region. This figure explains the observed line profile in ASASSN-14il at early
phases ($<$ day 50).

Between day 50–80, the photosphere recedes into the interaction region. Due to
the asymmetric CSM geometry, both uninterrupted SN ejecta and post-shock
region are visible simultaneously (Figure 11b). The line profiles see
contributions from the SN ejecta, shock region, and distant uninterrupted CSM,
which manifests as broad, intermediate, and narrow-width Gaussian components,
respectively.

Figure 11c depicts the evolution of ASASSN-14il post day 80. The contributions
from the distant CSM, post-shock region, and the ejecta are still seen. Dust
formation is seen in the post-shock gas, an efficient location for dust
formation in SNe IIn. This dust preferentially obscures the receding SN ejecta
and shock, causing a deficit in the redside wing of line profiles. This effect
manifests as an overall blueshift of the broad component. Such blueshift is
observed in the late time ($>$ day 100) H$\alpha$ profile of ASASSN-14il as
well as many other SNe IIn (e.g. SNe 2005ip, 20006tf, 2010jl, ASASSN-15ua,
2015da).

Similar explosion/CSM geometry have been proposed for SNe 2010jl, 2014ab,
2015da, and 2017hcc (Smith et al., 2011; Bilinski et al., 2020; Smith et al.,
2023; Smith & Andrews, 2020) with different viewing angles. However, the
viewing angle in SN 2014ab and ASASSN-14il are proposed to be similar, which
is further emphasized by both having remarkably low polarization (Bilinski et
al., 2024).

##  8 Summary

  1. 1.

ASASSN-14il was observed with an extensive follow-up campaign spanning $\sim
480$ days in photometric and $\sim 327$ days in spectroscopy observations.

  2. 2.

ASASSN-14il is a very luminous SN IIn showing long-term interaction signatures
and a plateau in the optical light curves around the maximum. It peaks at
$\sim-20.3$ mag in r-band ($<-21$ mag in UV bands), comparable to SLSNe
2006tf, 2010jl, ASASSN-15ua. The lightcurve shape is very similar to SN
2015da, but the late time (after $\sim 250$ days) decline rate of $\sim$0.01
mag/day in optical bands is faster than SN 2015da but comparable to SNe
2006tf, 2010jl, ASASSN-15ua. The slow reward evolution of the intrinsic B-V
resembles the typical SNe IIn.

  3. 3.

The multi-band lightcurve modeling of ASASSN-14il indicates a CSM-driven
explosion. Both wind/shell CSM model generates our observed lightcurves with a
CSM mass between 4.7–9.1 M⊙ and the CSM shell/wind being expelled about 1-5
years before the explosion.

  4. 4.

Spectroscopically, ASASSN-14il shows long-term predominant interaction
signatures with narrow H and He lines on top of otherwise featureless spectra
similar to other SLSNe IIn 2006gy, 2010jl, and 2015da.

  5. 5.

The H$\alpha$ profile of ASASSN-14il can be described by the following
components- (i) An unresolved component from the pre-shock ionized CSM that
stays consistent throughout the evolution, (ii) An intermediate width
component initially dominated by Lorentzian electron-scattering wings but
later (day $>50$) show Gaussian profile (FWHM $\sim$ 1750 km s-1)
representative of the post-shock region. (iii) A broad Gaussian emission
component (FWHM $\sim$ 7000 km s-1) representative of the uninterrupted
ejecta, which becomes visible (day $>50$).

  6. 6.

The emergence of the ejecta component when the lightcurves are near peak
luminosity indicates asymmetry in the CSM structure.

  7. 7.

The blueshift of H$\alpha$ profile at late phases (day $>100$) indicates dust
formation in the post-shock CSM/ejecta.

  8. 8.

The mass-loss rate of 2-7 M${}_{\odot}\mathrm{yr}^{-1}$ indicates that the
progenitor star is most likely an LBV that underwent an eruptive phase.

## Acknowledgements

This work uses data from the Las Cumbres Observatory Global Telescope network.
The LCO group is supported by NSF grant AST-1911151. This work uses Swift/UVOT
data reduced by P. J. Brown and released in the Swift Optical/Ultraviolet
Supernova Archive (SOUSA). SOUSA is supported by NASA’s Astrophysics Data
Analysis Program through grant NNX13AF35G. This research has made use of the
APASS database, located on the AAVSO website. Funding for APASS has been
provided by the Robert Martin Ayers Sciences Fund. This work is based (in
part) on observations collected at the European Organisation for Astronomical
Research in the Southern Hemisphere, Chile as part of PESSTO, (the Public ESO
Spectroscopic Survey for Transient Objects Survey) ESO program 188.D-3003,
191.D-0935, 197.D-1075. B.A. acknowledges the Council of Scientific $\&$
Industrial Research (CSIR) fellowship award (09/948(0005)/2020-EMR-I) for this
work.

## References

  * Agnoletto (2009) Agnoletto, I. 2009, in American Institute of Physics Conference Series, Vol. 1111, American Institute of Physics Conference Series, ed. G. Giobbi, A. Tornambe, G. Raimondo, M. Limongi, L. A. Antonelli, N. Menci, & E. Brocato, 430–434, doi: 10.1063/1.3141587
  * Bak Nielsen et al. (2018) Bak Nielsen, A.-S., Hjorth, J., & Gall, C. 2018, A&A, 611, A67, doi: 10.1051/0004-6361/201629904
  * Balberg & Loeb (2011) Balberg, S., & Loeb, A. 2011, MNRAS, 414, 1715, doi: 10.1111/j.1365-2966.2011.18505.x
  * Becker (2015) Becker, A. 2015, HOTPANTS: High Order Transform of PSF ANd Template Subtraction, Astrophysics Source Code Library, record ascl:1504.004.  http://ascl.net/1504.004
  * Benetti et al. (2006) Benetti, S., Cappellaro, E., Turatto, M., et al. 2006, ApJ, 653, L129, doi: 10.1086/510667
  * Bilinski et al. (2024) Bilinski, C., Smith, N., Williams, G. G., et al. 2024, MNRAS, 529, 1104, doi: 10.1093/mnras/stae380
  * Bilinski et al. (2020) —. 2020, MNRAS, 498, 3835, doi: 10.1093/mnras/staa2617
  * Breeveld et al. (2011) Breeveld, A. A., Landsman, W., Holland, S. T., et al. 2011, in American Institute of Physics Conference Series, Vol. 1358, Gamma Ray Bursts 2010, ed. J. E. McEnery, J. L. Racusin, & N. Gehrels, 373–376, doi: 10.1063/1.3621807
  * Brimacombe et al. (2014) Brimacombe, J., Holoien, T. W. S., Stanek, K. Z., et al. 2014, The Astronomer’s Telegram, 6525, 1
  * Brown et al. (2014) Brown, P. J., Breeveld, A. A., Holland, S., Kuin, P., & Pritchard, T. 2014, Ap&SS, 354, 89, doi: 10.1007/s10509-014-2059-8
  * Brown et al. (2009) Brown, P. J., Holland, S. T., Immler, S., et al. 2009, AJ, 137, 4517, doi: 10.1088/0004-6256/137/5/4517
  * Chatzopoulos et al. (2012) Chatzopoulos, E., Wheeler, J. C., & Vinko, J. 2012, ApJ, 746, 121, doi: 10.1088/0004-637X/746/2/121
  * Chatzopoulos et al. (2013) Chatzopoulos, E., Wheeler, J. C., Vinko, J., Horvath, Z. L., & Nagy, A. 2013, ApJ, 773, 76, doi: 10.1088/0004-637X/773/1/76
  * Chevalier (1982) Chevalier, R. A. 1982, ApJ, 258, 790, doi: 10.1086/160126
  * Chevalier (1998) —. 1998, ApJ, 499, 810, doi: 10.1086/305676
  * Chevalier & Fransson (1994) Chevalier, R. A., & Fransson, C. 1994, ApJ, 420, 268, doi: 10.1086/173557
  * Childress et al. (2014) Childress, M., Scalzo, R., Yuan, F., et al. 2014, The Astronomer’s Telegram, 6536, 1
  * Childress et al. (2016) Childress, M. J., Tucker, B. E., Yuan, F., et al. 2016, PASA, 33, e055, doi: 10.1017/pasa.2016.47
  * Chugai (2001) Chugai, N. N. 2001, MNRAS, 326, 1448, doi: 10.1111/j.1365-2966.2001.04717.x
  * Chugai & Danziger (1994) Chugai, N. N., & Danziger, I. J. 1994, MNRAS, 268, 173, doi: 10.1093/mnras/268.1.173
  * Chugai et al. (2004) Chugai, N. N., Blinnikov, S. I., Cumming, R. J., et al. 2004, MNRAS, 352, 1213, doi: 10.1111/j.1365-2966.2004.08011.x
  * Dessart et al. (2015) Dessart, L., Audit, E., & Hillier, D. J. 2015, MNRAS, 449, 4304, doi: 10.1093/mnras/stv609
  * Dessart et al. (2009) Dessart, L., Hillier, D. J., Gezari, S., Basa, S., & Matheson, T. 2009, MNRAS, 394, 21, doi: 10.1111/j.1365-2966.2008.14042.x
  * Di Carlo et al. (2002) Di Carlo, E., Massi, F., Valentini, G., et al. 2002, ApJ, 573, 144, doi: 10.1086/340496
  * Dickinson et al. (2023) Dickinson, D., Smith, N., Andrews, J. E., et al. 2023, arXiv e-prints, arXiv:2302.04958, doi: 10.48550/arXiv.2302.04958
  * Dickinson (2021) Dickinson, D. A. 2021, AN OPTICAL STUDY OF THE SUPERLUMINOUS TYPE IIN SUPERNOVA ASASSN-14IL.  http://hdl.handle.net/10150/666597
  * Filippenko (1997) Filippenko, A. V. 1997, ARA&A, 35, 309, doi: 10.1146/annurev.astro.35.1.309
  * Fox et al. (2009) Fox, O., Skrutskie, M. F., Chevalier, R. A., et al. 2009, ApJ, 691, 650, doi: 10.1088/0004-637X/691/1/650
  * Fox et al. (2015) Fox, O. D., Silverman, J. M., Filippenko, A. V., et al. 2015, MNRAS, 447, 772, doi: 10.1093/mnras/stu2435
  * Fransson et al. (2002) Fransson, C., Chevalier, R. A., Filippenko, A. V., et al. 2002, ApJ, 572, 350, doi: 10.1086/340295
  * Fransson et al. (2014) Fransson, C., Ergon, M., Challis, P. J., et al. 2014, ApJ, 797, 118, doi: 10.1088/0004-637X/797/2/118
  * Gal-Yam (2012) Gal-Yam, A. 2012, Science, 337, 927, doi: 10.1126/science.1203601
  * Gal-Yam (2019) —. 2019, ARA&A, 57, 305, doi: 10.1146/annurev-astro-081817-051819
  * Gal-Yam & Leonard (2009) Gal-Yam, A., & Leonard, D. C. 2009, Nature, 458, 865, doi: 10.1038/nature07934
  * Gall et al. (2014) Gall, C., Hjorth, J., Watson, D., et al. 2014, Nature, 511, 326, doi: 10.1038/nature13558
  * Gangopadhyay et al. (2020) Gangopadhyay, A., Turatto, M., Benetti, S., et al. 2020, MNRAS, 499, 129, doi: 10.1093/mnras/staa2606
  * Gehrels et al. (2004) Gehrels, N., Chincarini, G., Giommi, P., et al. 2004, ApJ, 611, 1005, doi: 10.1086/422091
  * Graham et al. (2014) Graham, M. L., Sand, D. J., Valenti, S., et al. 2014, ApJ, 787, 163, doi: 10.1088/0004-637X/787/2/163
  * Guillochon et al. (2018) Guillochon, J., Nicholl, M., Villar, V. A., et al. 2018, ApJS, 236, 6, doi: 10.3847/1538-4365/aab761
  * Hosseinzadeh et al. (2022) Hosseinzadeh, G., Berger, E., Metzger, B. D., et al. 2022, ApJ, 933, 14, doi: 10.3847/1538-4357/ac67dd
  * Hosseinzadeh & Gomez (2022) Hosseinzadeh, G., & Gomez, S. 2022, Light Curve Fitting, v0.6.0, Zenodo, doi: 10.5281/zenodo.6519623
  * Ishii et al. (2024) Ishii, A. T., Takei, Y., Tsuna, D., Shigeyama, T., & Takahashi, K. 2024, ApJ, 961, 47, doi: 10.3847/1538-4357/ad072b
  * Jencson et al. (2016) Jencson, J. E., Prieto, J. L., Kochanek, C. S., et al. 2016, MNRAS, 456, 2622, doi: 10.1093/mnras/stv2795
  * Jiang et al. (2020) Jiang, B., Jiang, S., & Ashley Villar, V. 2020, Research Notes of the American Astronomical Society, 4, 16, doi: 10.3847/2515-5172/ab7128
  * Jones et al. (2004) Jones, D. H., Saunders, W., Colless, M., et al. 2004, MNRAS, 355, 747, doi: 10.1111/j.1365-2966.2004.08353.x
  * Jones et al. (2009) Jones, D. H., Read, M. A., Saunders, W., et al. 2009, MNRAS, 399, 683, doi: 10.1111/j.1365-2966.2009.15338.x
  * Kankare et al. (2012) Kankare, E., Ergon, M., Bufano, F., et al. 2012, MNRAS, 424, 855, doi: 10.1111/j.1365-2966.2012.21224.x
  * Katsuda et al. (2016) Katsuda, S., Maeda, K., Bamba, A., et al. 2016, ApJ, 832, 194, doi: 10.3847/0004-637X/832/2/194
  * Kiewe et al. (2012) Kiewe, M., Gal-Yam, A., Arcavi, I., et al. 2012, ApJ, 744, 10, doi: 10.1088/0004-637X/744/1/10
  * Kochanek et al. (2011) Kochanek, C. S., Szczygiel, D. M., & Stanek, K. Z. 2011, ApJ, 737, 76, doi: 10.1088/0004-637X/737/2/76
  * Li et al. (2011) Li, W., Leaman, J., Chornock, R., et al. 2011, MNRAS, 412, 1441, doi: 10.1111/j.1365-2966.2011.18160.x
  * Maeda et al. (2013) Maeda, K., Nozawa, T., Sahu, D. K., et al. 2013, ApJ, 776, 5, doi: 10.1088/0004-637X/776/1/5
  * Mattila et al. (2008) Mattila, S., Meikle, W. P. S., Lundqvist, P., et al. 2008, MNRAS, 389, 141, doi: 10.1111/j.1365-2966.2008.13516.x
  * Matzner & McKee (1999) Matzner, C. D., & McKee, C. F. 1999, ApJ, 510, 379, doi: 10.1086/306571
  * Mauerhan et al. (2013) Mauerhan, J. C., Smith, N., Silverman, J. M., et al. 2013, MNRAS, 431, 2599, doi: 10.1093/mnras/stt360
  * Moriya et al. (2013) Moriya, T. J., Maeda, K., Taddia, F., et al. 2013, MNRAS, 435, 1520, doi: 10.1093/mnras/stt1392
  * Moriya et al. (2014) —. 2014, MNRAS, 439, 2917, doi: 10.1093/mnras/stu163
  * Nicholl (2018) Nicholl, M. 2018, Research Notes of the American Astronomical Society, 2, 230, doi: 10.3847/2515-5172/aaf799
  * Nicholl et al. (2020) Nicholl, M., Blanchard, P. K., Berger, E., et al. 2020, Nature Astronomy, 4, 893, doi: 10.1038/s41550-020-1066-7
  * Nyholm et al. (2020) Nyholm, A., Sollerman, J., Tartaglia, L., et al. 2020, A&A, 637, A73, doi: 10.1051/0004-6361/201936097
  * Ofek et al. (2014) Ofek, E. O., Zoglauer, A., Boggs, S. E., et al. 2014, ApJ, 781, 42, doi: 10.1088/0004-637X/781/1/42
  * Poznanski et al. (2012) Poznanski, D., Prochaska, J. X., & Bloom, J. S. 2012, MNRAS, 426, 1465, doi: 10.1111/j.1365-2966.2012.21796.x
  * Pozzo et al. (2004) Pozzo, M., Meikle, W. P. S., Fassia, A., et al. 2004, MNRAS, 352, 457, doi: 10.1111/j.1365-2966.2004.07951.x
  * Roming et al. (2005) Roming, P. W. A., Kennedy, T. E., Mason, K. O., et al. 2005, Space Sci. Rev., 120, 95, doi: 10.1007/s11214-005-5095-4
  * Schlafly & Finkbeiner (2011) Schlafly, E. F., & Finkbeiner, D. P. 2011, ApJ, 737, 103, doi: 10.1088/0004-637X/737/2/103
  * Schlegel (1990) Schlegel, E. M. 1990, MNRAS, 244, 269
  * Silverman et al. (2013) Silverman, J. M., Nugent, P. E., Gal-Yam, A., et al. 2013, ApJ, 772, 125, doi: 10.1088/0004-637X/772/2/125
  * Smartt et al. (2015) Smartt, S. J., Valenti, S., Fraser, M., et al. 2015, A&A, 579, A40, doi: 10.1051/0004-6361/201425237
  * Smith (2014) Smith, N. 2014, ARA&A, 52, 487, doi: 10.1146/annurev-astro-081913-040025
  * Smith (2017) —. 2017, in Handbook of Supernovae, ed. A. W. Alsabti & P. Murdin (Springer), 403, doi: 10.1007/978-3-319-21846-5_38
  * Smith & Andrews (2020) Smith, N., & Andrews, J. E. 2020, MNRAS, 499, 3544, doi: 10.1093/mnras/staa3047
  * Smith et al. (2023) Smith, N., Andrews, J. E., Milne, P., et al. 2023, arXiv e-prints, arXiv:2312.00253, doi: 10.48550/arXiv.2312.00253
  * Smith et al. (2008) Smith, N., Chornock, R., Li, W., et al. 2008, ApJ, 686, 467, doi: 10.1086/591021
  * Smith et al. (2009a) Smith, N., Hinkle, K. H., & Ryde, N. 2009a, AJ, 137, 3558, doi: 10.1088/0004-6256/137/3/3558
  * Smith & McCray (2007) Smith, N., & McCray, R. 2007, ApJ, 671, L17, doi: 10.1086/524681
  * Smith & Owocki (2006) Smith, N., & Owocki, S. P. 2006, ApJ, 645, L45, doi: 10.1086/506523
  * Smith et al. (2012) Smith, N., Silverman, J. M., Filippenko, A. V., et al. 2012, AJ, 143, 17, doi: 10.1088/0004-6256/143/1/17
  * Smith et al. (2007) Smith, N., Li, W., Foley, R. J., et al. 2007, ApJ, 666, 1116, doi: 10.1086/519949
  * Smith et al. (2009b) Smith, N., Silverman, J. M., Chornock, R., et al. 2009b, ApJ, 695, 1334, doi: 10.1088/0004-637X/695/2/1334
  * Smith et al. (2010) Smith, N., Miller, A., Li, W., et al. 2010, in American Astronomical Society Meeting Abstracts, Vol. 215, American Astronomical Society Meeting Abstracts #215, 430.16
  * Smith et al. (2011) Smith, N., Li, W., Miller, A. A., et al. 2011, ApJ, 732, 63, doi: 10.1088/0004-637X/732/2/63
  * Speagle (2020) Speagle, J. S. 2020, MNRAS, 493, 3132, doi: 10.1093/mnras/staa278
  * Stritzinger et al. (2012) Stritzinger, M., Taddia, F., Fransson, C., et al. 2012, ApJ, 756, 173, doi: 10.1088/0004-637X/756/2/173
  * Suzuki et al. (2021) Suzuki, A., Nicholl, M., Moriya, T. J., & Takiwaki, T. 2021, ApJ, 908, 99, doi: 10.3847/1538-4357/abd6ce
  * Svirski et al. (2012) Svirski, G., Nakar, E., & Sari, R. 2012, ApJ, 759, 108, doi: 10.1088/0004-637X/759/2/108
  * Taddia et al. (2013) Taddia, F., Stritzinger, M. D., Sollerman, J., et al. 2013, A&A, 555, A10, doi: 10.1051/0004-6361/201321180
  * Tartaglia et al. (2020) Tartaglia, L., Pastorello, A., Sollerman, J., et al. 2020, A&A, 635, A39, doi: 10.1051/0004-6361/201936553
  * Valenti et al. (2014) Valenti, S., Sand, D., Pastorello, A., et al. 2014, MNRAS, 438, L101, doi: 10.1093/mnrasl/slt171
  * Valenti et al. (2016) Valenti, S., Howell, D. A., Stritzinger, M. D., et al. 2016, MNRAS, 459, 3939, doi: 10.1093/mnras/stw870
  * Vigna-Gómez et al. (2019) Vigna-Gómez, A., Justham, S., Mandel, I., de Mink, S. E., & Podsiadlowski, P. 2019, ApJ, 876, L29, doi: 10.3847/2041-8213/ab1bdf
  * Vink (2018) Vink, J. S. 2018, A&A, 619, A54, doi: 10.1051/0004-6361/201833352
  * Woosley (2017) Woosley, S. E. 2017, ApJ, 836, 244, doi: 10.3847/1538-4357/836/2/244
  * Woosley et al. (2007) Woosley, S. E., Blinnikov, S., & Heger, A. 2007, Nature, 450, 390, doi: 10.1038/nature06333
  * Yaron & Gal-Yam (2012) Yaron, O., & Gal-Yam, A. 2012, PASP, 124, 668, doi: 10.1086/666656
  * Zhang et al. (2012) Zhang, T., Wang, X., Wu, C., et al. 2012, AJ, 144, 131, doi: 10.1088/0004-6256/144/5/131

##  Appendix A Observation log

Table 7: Log of photomteric observations of ASASSN-14il obtained from LCOGT network of telescopes.  Date | Phase | B | V | g | r | i | telescope
---|---|---|---|---|---|---|---
(yyyy-mm-dd) | (d) | (mag) | (mag) | (mag) | (mag) | (mag) |
2014-10-03 | 7.8 | $16.15\pm 0.03$ | $16.25\pm 0.04$ | $15.97\pm 0.02$ | $16.10\pm 0.04$ | $16.14\pm 0.01$ | 1m0-12
2014-10-05 | 9.5 | $16.15\pm 0.02$ | $16.07\pm 0.01$ | $15.87\pm 0.04$ | $15.91\pm 0.02$ | $16.07\pm 0.01$ | 1m0-11
2014-10-07 | 11.3 | $16.10\pm 0.08$ | $15.71\pm 0.06$ | $15.75\pm 0.05$ | $15.84\pm 0.11$ | - | 1m0-08
2014-10-09 | 13.3 | - | $15.58\pm 0.16$ | $15.57\pm 0.05$ | $15.54\pm 0.15$ | $15.70\pm 0.20$ | 1m0-08
2014-10-09 | 13.7 | $15.87\pm 0.02$ | $15.74\pm 0.03$ | $15.50\pm 0.02$ | $15.52\pm 0.04$ | $15.65\pm 0.05$ | 1m0-11
2014-10-11 | 15.5 | $15.84\pm 0.06$ | $15.63\pm 0.01$ | $15.42\pm 0.01$ | $15.52\pm 0.03$ | $15.68\pm 0.05$ | 1m0-03
2014-10-13 | 17.9 | $15.76\pm 0.04$ | $15.55\pm 0.01$ | $15.44\pm 0.01$ | $15.45\pm 0.02$ | $15.53\pm 0.01$ | 1m0-10
2014-10-15 | 19.3 | $15.74\pm 0.05$ | $15.57\pm 0.03$ | $15.37\pm 0.02$ | $15.42\pm 0.02$ | $15.56\pm 0.02$ | 1m0-08
2014-10-17 | 22.0 | $15.74\pm 0.03$ | $15.44\pm 0.02$ | $15.39\pm 0.01$ | $15.37\pm 0.04$ | $15.44\pm 0.04$ | 1m0-13
2014-10-18 | 22.2 | $15.63\pm 0.01$ | $15.38\pm 0.03$ | $15.23\pm 0.01$ | $15.33\pm 0.01$ | $15.39\pm 0.05$ | 1m0-08
2014-10-19 | 23.9 | $15.68\pm 0.02$ | - | - | - | - | 1m0-12
2014-10-20 | 24.5 | $15.54\pm 0.06$ | $15.40\pm 0.04$ | $15.27\pm 0.02$ | $15.22\pm 0.02$ | $15.28\pm 0.03$ | 1m0-03
2014-10-20 | 24.4 | $15.54\pm 0.02$ | $15.30\pm 0.03$ | $15.21\pm 0.01$ | $15.21\pm 0.01$ | $15.27\pm 0.02$ | 1m0-11
2014-10-21 | 25.9 | $15.70\pm 0.02$ | $15.37\pm 0.02$ | $15.30\pm 0.01$ | $15.25\pm 0.03$ | $15.31\pm 0.02$ | 1m0-10
2014-10-22 | 26.5 | $15.59\pm 0.02$ | $15.30\pm 0.04$ | $15.18\pm 0.03$ | - | $15.31\pm 0.10$ | 1m0-03
2014-10-26 | 30.5 | $15.65\pm 0.02$ | $15.32\pm 0.07$ | $15.22\pm 0.02$ | $15.13\pm 0.02$ | $15.17\pm 0.05$ | 1m0-11
2014-10-29 | 34.0 | - | $15.33\pm 0.02$ | $15.24\pm 0.01$ | $15.14\pm 0.03$ | $15.19\pm 0.05$ | 1m0-12
2014-11-02 | 37.7 | - | $15.32\pm 0.09$ | - | - | - | 1m0-03
2014-11-06 | 41.4 | - | $15.15\pm 0.03$ | $15.23\pm 0.02$ | $15.04\pm 0.03$ | - | 1m0-11
2014-11-10 | 45.2 | - | $15.20\pm 0.04$ | $15.32\pm 0.04$ | $15.10\pm 0.04$ | $15.09\pm 0.03$ | 1m0-08
2014-11-13 | 49.0 | $15.73\pm 0.03$ | $15.22\pm 0.02$ | $15.23\pm 0.03$ | $15.06\pm 0.03$ | $15.06\pm 0.02$ | 1m0-12
2014-11-17 | 52.8 | $15.62\pm 0.02$ | $15.26\pm 0.02$ | $15.34\pm 0.01$ | $15.11\pm 0.04$ | $15.02\pm 0.03$ | 1m0-12
2014-11-21 | 56.6 | - | $15.26\pm 0.04$ | $15.27\pm 0.02$ | $15.04\pm 0.03$ | $14.98\pm 0.03$ | 1m0-03
2014-11-26 | 61.9 | $15.78\pm 0.02$ | - | $15.23\pm 0.02$ | $15.04\pm 0.02$ | $14.97\pm 0.02$ | 1m0-10
2014-11-26 | 62.0 | - | - | $15.31\pm 0.03$ | $15.05\pm 0.01$ | $14.97\pm 0.04$ | 1m0-13
2014-11-27 | 62.2 | $15.76\pm 0.03$ | $15.29\pm 0.04$ | $15.31\pm 0.01$ | $15.08\pm 0.02$ | $15.03\pm 0.02$ | 1m0-05
2014-11-27 | 62.0 | $15.67\pm 0.01$ | $15.18\pm 0.03$ | $15.26\pm 0.01$ | $15.04\pm 0.04$ | $15.00\pm 0.04$ | 1m0-08
2014-11-28 | 63.8 | $15.89\pm 0.04$ | $15.30\pm 0.02$ | $15.39\pm 0.01$ | $15.08\pm 0.03$ | $14.98\pm 0.05$ | 1m0-12
2014-12-02 | 67.4 | - | - | - | $15.09\pm 0.05$ | $15.10\pm 0.04$ | 1m0-03
2014-12-04 | 69.8 | $15.92\pm 0.02$ | $15.27\pm 0.03$ | $15.43\pm 0.02$ | $15.06\pm 0.02$ | $15.03\pm 0.04$ | 1m0-10
2014-12-04 | 69.9 | $15.81\pm 0.01$ | $15.27\pm 0.02$ | $15.38\pm 0.01$ | $15.07\pm 0.01$ | $14.92\pm 0.04$ | 1m0-12
2014-12-09 | 74.2 | $15.84\pm 0.02$ | $15.23\pm 0.03$ | $15.32\pm 0.02$ | $14.99\pm 0.03$ | $15.02\pm 0.01$ | 1m0-05
2014-12-12 | 77.8 | $15.93\pm 0.03$ | $15.17\pm 0.03$ | $15.48\pm 0.02$ | - | - | 1m0-12
2014-12-12 | 77.9 | $15.88\pm 0.03$ | $15.15\pm 0.06$ | $15.48\pm 0.05$ | $15.05\pm 0.03$ | $14.92\pm 0.01$ | 1m0-13
2014-12-16 | 81.8 | $15.97\pm 0.03$ | $15.28\pm 0.02$ | $15.34\pm 0.03$ | $15.05\pm 0.02$ | $14.89\pm 0.03$ | 1m0-12
2014-12-20 | 85.8 | $15.91\pm 0.04$ | $15.27\pm 0.01$ | $15.41\pm 0.02$ | $15.00\pm 0.05$ | $14.87\pm 0.03$ | 1m0-12
2014-12-24 | 89.5 | $15.72\pm 0.02$ | $15.10\pm 0.02$ | $15.22\pm 0.02$ | $14.89\pm 0.02$ | $14.83\pm 0.04$ | 1m0-03
2014-12-29 | 94.5 | $15.82\pm 0.01$ | $15.23\pm 0.10$ | $15.41\pm 0.02$ | - | - | 1m0-11
2015-01-02 | 98.5 | - | - | - | $14.83\pm 0.03$ | $14.78\pm 0.03$ | 1m0-11
2015-01-06 | 102.4 | $15.87\pm 0.03$ | $15.27\pm 0.01$ | $15.32\pm 0.01$ | $14.93\pm 0.02$ | $14.79\pm 0.05$ | 1m0-03
2015-01-11 | 107.8 | $16.07\pm 0.01$ | $15.31\pm 0.02$ | $15.54\pm 0.02$ | $15.04\pm 0.01$ | $14.89\pm 0.02$ | 1m0-13
2015-01-15 | 111.8 | $16.19\pm 0.04$ | $15.33\pm 0.01$ | $15.52\pm 0.02$ | - | - | 1m0-13
2015-01-16 | 112.1 | $15.98\pm 0.02$ | $15.33\pm 0.03$ | $15.48\pm 0.01$ | $15.05\pm 0.02$ | $14.96\pm 0.02$ | 1m0-05
2015-01-19 | 115.8 | $16.02\pm 0.02$ | $15.37\pm 0.04$ | $15.51\pm 0.02$ | $15.04\pm 0.02$ | $14.89\pm 0.05$ | 1m0-10
2015-01-28 | 124.0 | $16.24\pm 0.05$ | $15.58\pm 0.04$ | $15.63\pm 0.02$ | $15.14\pm 0.01$ | $15.08\pm 0.02$ | 1m0-05
2015-02-06 | 133.4 | $16.09\pm 0.08$ | $15.37\pm 0.01$ | $15.59\pm 0.01$ | $15.11\pm 0.02$ | $14.99\pm 0.06$ | 1m0-03
2015-06-19 | 266.2 | - | $17.09\pm 0.08$ | $17.33\pm 0.02$ | $16.53\pm 0.02$ | $16.44\pm 0.08$ | 1m0-12
2015-06-24 | 271.4 | - | - | $17.28\pm 0.02$ | $16.85\pm 0.09$ | $16.77\pm 0.02$ | 1m0-05
2015-06-28 | 275.8 | $17.97\pm 0.03$ | $16.97\pm 0.03$ | $17.44\pm 0.03$ | $16.59\pm 0.04$ | - | 1m0-11
2015-07-02 | 279.2 | $17.71\pm 0.15$ | $17.22\pm 0.04$ | $17.45\pm 0.07$ | $16.71\pm 0.01$ | $16.92\pm 0.02$ | 1m0-12
2015-07-02 | 279.2 | $17.79\pm 0.17$ | $17.26\pm 0.04$ | $17.24\pm 0.09$ | $16.61\pm 0.03$ | $16.76\pm 0.13$ | 1m0-13
2015-07-06 | 283.0 | $17.49\pm 0.20$ | - | $17.77\pm 0.11$ | $16.50\pm 0.19$ | $16.76\pm 0.03$ | 1m0-12
2015-07-09 | 286.7 | - | $17.40\pm 0.08$ | - | $16.53\pm 0.05$ | $16.76\pm 0.22$ | 1m0-11
2015-07-14 | 291.4 | $18.12\pm 0.05$ | $17.42\pm 0.09$ | $17.52\pm 0.03$ | $16.75\pm 0.04$ | $16.86\pm 0.02$ | 1m0-05
2015-07-18 | 295.4 | $18.38\pm 0.22$ | $17.37\pm 0.04$ | $17.38\pm 0.08$ | $16.57\pm 0.06$ | $16.92\pm 0.08$ | 1m0-08
2015-07-23 | 300.4 | $18.38\pm 0.16$ | $17.44\pm 0.06$ | $17.55\pm 0.04$ | $16.63\pm 0.06$ | $16.96\pm 0.01$ | 1m0-05
2015-07-26 | 303.7 | - | $17.31\pm 0.02$ | $17.60\pm 0.07$ | $16.63\pm 0.01$ | $16.84\pm 0.04$ | 1m0-03
2015-07-30 | 307.4 | $18.41\pm 0.05$ | $17.53\pm 0.13$ | - | $16.91\pm 0.02$ | $17.04\pm 0.09$ | 1m0-05
2015-08-25 | 333.4 | - | $17.82\pm 0.04$ | $17.81\pm 0.03$ | $16.99\pm 0.04$ | $17.11\pm 0.09$ | 1m0-08
2015-09-04 | 343.9 | $18.60\pm 0.13$ | $17.60\pm 0.12$ | $17.89\pm 0.02$ | $17.18\pm 0.03$ | $17.26\pm 0.08$ | 1m0-10
2015-09-05 | 344.1 | $18.45\pm 0.05$ | $17.92\pm 0.03$ | - | - | - | 1m0-10
2015-09-05 | 344.1 | - | $17.68\pm 0.30$ | $17.96\pm 0.12$ | $17.18\pm 0.07$ | $17.63\pm 0.04$ | 1m0-12
2015-09-05 | 344.0 | - | $17.90\pm 0.07$ | $18.05\pm 0.04$ | $17.04\pm 0.09$ | - | 1m0-13
2015-09-15 | 354.8 | $18.65\pm 0.30$ | $17.76\pm 0.13$ | $17.97\pm 0.04$ | - | $17.42\pm 0.13$ | 1m0-12
2015-09-18 | 357.3 | $18.89\pm 0.05$ | $18.26\pm 0.16$ | $18.13\pm 0.09$ | $17.27\pm 0.01$ | $17.56\pm 0.05$ | 1m0-05
2015-09-28 | 368.0 | - | - | $18.16\pm 0.08$ | - | - | 1m0-10
2015-09-29 | 368.0 | - | - | - | $17.28\pm 0.11$ | - | 1m0-10
2015-10-14 | 383.2 | - | $18.50\pm 0.06$ | $18.30\pm 0.14$ | $17.50\pm 0.02$ | $18.02\pm 0.12$ | 1m0-08
2015-10-27 | 396.2 | - | - | $18.53\pm 0.21$ | $17.64\pm 0.11$ | - | 1m0-05
2015-11-06 | 406.9 | $19.44\pm 0.12$ | - | $18.87\pm 0.03$ | $18.05\pm 0.01$ | $17.80\pm 0.19$ | 1m0-10
2015-11-17 | 417.6 | $19.49\pm 0.40$ | $18.80\pm 0.32$ | $19.18\pm 0.12$ | - | $18.06\pm 0.11$ | 1m0-03
2015-11-19 | 419.6 | - | $18.75\pm 0.03$ | - | - | $18.09\pm 0.23$ | 1m0-03
2015-11-21 | 421.2 | - | - | $18.88\pm 0.15$ | $18.01\pm 0.19$ | - | 1m0-05
2015-12-14 | 444.1 | $20.40\pm 0.23$ | - | $19.15\pm 0.11$ | - | $18.57\pm 0.09$ | 1m0-05
2015-12-25 | 455.1 | - | - | - | $18.72\pm 0.29$ | $18.47\pm 0.17$ | 1m0-05
2016-01-07 | 468.4 | $20.46\pm 0.18$ | - | - | $18.78\pm 0.18$ | - | 1m0-03
2016-01-07 | 468.1 | - | - | $19.59\pm 0.17$ | $18.40\pm 0.08$ | - | 1m0-05
2016-01-07 | 468.8 | - | - | - | $18.66\pm 0.23$ | - | 1m0-10
2016-01-19 | 480.1 | - | - | - | $18.74\pm 0.08$ | - | 1m0-05
Table 8: Log of Swift UV and optical photometric observations of ASASSN-14il. Phase | MJD | UVW1 | UVW2 | UVM2 | U | B | V | Instrument
---|---|---|---|---|---|---|---|---
(d) |  | (mag) | (mag) | (mag) | (mag) | (mag) | (mag) |
7.90 | 56933.90 | $14.70\pm 0.08$ | $14.72\pm 0.07$ | $14.69\pm 0.08$ | $14.81\pm 0.07$ | $16.08\pm 0.11$ | $16.08\pm 0.16$ | UVOT
9.99 | 56935.99 | $14.46\pm 0.07$ | $14.57\pm 0.07$ | $14.45\pm 0.06$ | $14.68\pm 0.07$ | $16.03\pm 0.10$ | $15.88\pm 0.14$ | UVOT
11.84 | 56937.84 | $14.41\pm 0.06$ | $14.54\pm 0.06$ | $14.42\pm 0.06$ | $14.54\pm 0.06$ | $15.83\pm 0.09$ | $15.77\pm 0.11$ | UVOT
13.83 | 56939.83 | $14.36\pm 0.07$ | $14.50\pm 0.07$ | $14.39\pm 0.06$ | $14.44\pm 0.06$ | $15.67\pm 0.09$ | $15.64\pm 0.12$ | UVOT
15.43 | 56941.43 | $14.29\pm 0.06$ | $14.49\pm 0.07$ | $14.37\pm 0.06$ | $14.40\pm 0.06$ | $15.55\pm 0.08$ | $15.55\pm 0.11$ | UVOT
16.88 | 56942.88 | $14.29\pm 0.06$ | $14.58\pm 0.07$ | - | $14.37\pm 0.06$ | $15.64\pm 0.09$ | - | UVOT
16.94 | 56942.94 | - | $14.58\pm 0.07$ | $14.53\pm 0.08$ | - | - | $15.66\pm 0.14$ | UVOT
17.83 | 56943.83 | $14.35\pm 0.06$ | $14.61\pm 0.07$ | $14.50\pm 0.07$ | $14.35\pm 0.06$ | $15.57\pm 0.08$ | $15.68\pm 0.12$ | UVOT
18.69 | 56944.69 | - | $14.63\pm 0.07$ | - | $14.32\pm 0.06$ | - | - | UVOT
18.87 | 56944.87 | $14.28\pm 0.06$ | $14.59\pm 0.07$ | $14.55\pm 0.07$ | $14.33\pm 0.06$ | $15.60\pm 0.09$ | $15.50\pm 0.11$ | UVOT
21.37 | 56947.37 | $14.39\pm 0.06$ | $14.66\pm 0.07$ | $14.57\pm 0.07$ | $14.34\pm 0.06$ | $15.55\pm 0.09$ | $15.25\pm 0.10$ | UVOT
23.25 | 56949.25 | $14.38\pm 0.06$ | $14.70\pm 0.07$ | $14.59\pm 0.07$ | $14.33\pm 0.06$ | $15.40\pm 0.08$ | $15.27\pm 0.10$ | UVOT
25.37 | 56951.37 | - | - | $14.58\pm 0.08$ | - | - | $15.25\pm 0.12$ | UVOT
25.67 | 56951.67 | $14.45\pm 0.07$ | $14.79\pm 0.08$ | - | $14.31\pm 0.06$ | $15.42\pm 0.08$ | - | UVOT
29.42 | 56955.42 | $14.55\pm 0.07$ | $14.92\pm 0.08$ | $14.78\pm 0.08$ | $14.38\pm 0.06$ | $15.40\pm 0.08$ | $15.25\pm 0.10$ | UVOT
31.34 | 56957.34 | $14.63\pm 0.07$ | $15.02\pm 0.08$ | $14.93\pm 0.08$ | $14.46\pm 0.06$ | $15.34\pm 0.07$ | $15.23\pm 0.10$ | UVOT
33.47 | 56959.47 | $14.77\pm 0.08$ | - | - | - | - | - | UVOT
35.92 | 56961.92 | $14.88\pm 0.08$ | $15.36\pm 0.10$ | $15.18\pm 0.09$ | $14.52\pm 0.06$ | $15.43\pm 0.08$ | $15.21\pm 0.09$ | UVOT
37.85 | 56963.85 | $15.00\pm 0.08$ | $15.41\pm 0.10$ | $15.27\pm 0.09$ | $14.60\pm 0.06$ | $15.41\pm 0.08$ | $15.23\pm 0.09$ | UVOT
40.78 | 56966.78 | $15.13\pm 0.09$ | $15.56\pm 0.11$ | $15.40\pm 0.09$ | $14.63\pm 0.07$ | $15.40\pm 0.08$ | $15.23\pm 0.09$ | UVOT
42.21 | 56968.21 | $15.14\pm 0.09$ | - | - | $14.64\pm 0.07$ | $15.47\pm 0.08$ | - | UVOT
46.23 | 56972.23 | $15.42\pm 0.10$ | $16.02\pm 0.14$ | $15.79\pm 0.11$ | $14.78\pm 0.07$ | $15.46\pm 0.08$ | $15.15\pm 0.09$ | UVOT
48.31 | 56974.31 | $15.45\pm 0.10$ | $16.05\pm 0.14$ | $15.87\pm 0.11$ | $14.85\pm 0.07$ | $15.57\pm 0.08$ | $15.18\pm 0.09$ | UVOT
Table 9: Log of spectroscopic observations of ASASSN-14il. Date | MJD | phase | Telescope | Range | Resolution
---|---|---|---|---|---
(yy-mm-dd) | (day) |  | (Å) | ($\lambda/\Delta\lambda$) |
2014-10-03 | 56933.55 | 7.6 | ANU+WiFeS | 3050 - 11000 | 3000
2014-10-03 | 56933.73 | 7.7 | LCO+FLOYDS | 3200 - 10000 | 400 - 700
2014-10-05 | 56935.63 | 9.6 | LCO+FLOYDS | 3200 - 10000 | 400 - 700
2014-10-07 | 56937.71 | 11.7 | LCO+FLOYDS | 3200 - 10000 | 400 - 700
2014-10-09 | 56939.50 | 13.5 | ANU+WiFeS | 3050 - 11000 | 3000
2014-10-10 | 56940.47 | 14.5 | LCO+FLOYDS | 3200 - 10000 | 400 - 700
2014-10-13 | 56943.64 | 17.6 | LCO+FLOYDS | 3200 - 10000 | 400 - 700
2014-10-14 | 56944.06 | 18.1 | ESO-NTT + SOFI | 9000 - 25000 | 600 - 2200
2014-10-18 | 56948.56 | 22.6 | ANU+WiFeS | 3050 - 11000 | 3000
2014-10-18 | 56948.67 | 22.7 | LCO+FLOYDS | 3200 - 10000 | 400 - 700
2014-10-22 | 56952.02 | 26.0 | ESO-NTT + EFOSC | 3050 - 11000 | -
2014-10-22 | 56952.47 | 26.5 | LCO+FLOYDS | 3200 - 10000 | 400 - 700
2014-10-24 | 56954.21 | 28.2 | ESO-NTT + SOFI | 9000 - 25000 | 600 - 2200
2014-10-27 | 56957.47 | 31.5 | ANU+WiFeS | 3050 - 11000 | 3000
2014-11-14 | 56975.16 | 49.2 | ESO-NTT + EFOSC | 3050 - 11000 | -
2014-11-15 | 56976.19 | 50.2 | ESO-NTT + SOFI | 9000 - 25000 | 600 - 2200
2014-12-16 | 57007.51 | 81.5 | LCO+FLOYDS | 3200 - 10000 | 400 - 700
2014-12-22 | 57013.08 | 87.1 | ESO-NTT + SOFI | 9000 - 25000 | 600 - 2200
2015-01-13 | 57035.06 | 109.1 | ESO-NTT + EFOSC | 3050 - 11000 | -
2015-01-20 | 57042.06 | 116.1 | ESO-NTT + SOFI | 9000 - 25000 | 600 - 2200
2015-01-27 | 57049.04 | 123.0 | ESO-NTT + EFOSC | 3050 - 11000 | -
2015-08-19 | 57253.26 | 327.3 | ESO-NTT + EFOSC | 3050 - 11000 | -

##  Appendix B MOSFiT Corner Plot

Figure 12: Corner plot showing the best fit parameters of the csm wind (blue)
and csm shell (gray) models for ASASSN-14il.