# Signatures of Circumstellar Interaction in the Type IIL Supernova
ASASSN-15oz

K. Azalee Bostroem, 1 Stefano Valenti,1 Assaf Horesh,2 Viktoriya Morozova,3 N.
Paul M. Kuin,4 Samuel Wyatt,5 Anders Jerkstrand,6 David J. Sand,5 Michael
Lundquist,5 Mathew Smith,7 Mark Sullivan,7 Griffin Hosseinzadeh,8 Iair
Arcavi,9 Emma Callis,10 Régis Cartier,11 Avishay Gal-Yam,12 Lluís Galbany,13
Claudia Gutiérrez,7 D. Andrew Howell,14, 15 Cosimo Inserra,17 Erkki Kankare,16
Kristhell Marisol López,18, 19 Curtis McCully,14, 15 Giuliano Pignata,20, 21
Anthony L. Piro,22 Ósmar Rodríguez,20, 21 Stephen J. Smartt,16 Kenneth W.
Smith,16 Ofer Yaron,12 David R. Young16
1Department of Physics, University of California, Davis, CA 95616, USA
2Racah Institute of Physics, The Hebrew University of Jerusalem, Jerusalem,
91904, Israel
3Department of Astrophysical Sciences, Princeton University, Princeton, NJ
08544, USA
4Mullard Space Science Laboratory, University College London, Holmbury St.
Mary, Dorking, Surrey RH5 6NT, UK
5Department of Astronomy/Steward Observatory, 933 North Cherry Avenue, Rm.
N204, Tucson, AZ 85721-0065, USA
6Max-Planck Institut für Astrophysik, Karl-Schwarzschild Str. 1, D-85748
Garching, Germany
7School of Physics and Astronomy, University of Southampton, Southampton, SO17
1BJ, UK
8Harvard-Smithsonian Center for Astrohysics, 60 Garden Street, Cambridge, MA
02138
9The Raymond and Beverly Sackler School of Physics and Astronomy, Tel Aviv
University, Tel Aviv 69978, Israel
10School of Physics, O’Brien Centre for Science North, University College
Dublin, Belfield, Dublin 4, Ireland
11Cerro Tololo Inter-American Observatory, National Optical Astronomy
Observatory, Casilla 603, La Serena, Chile
12Department of Particle Physics and Astrophysics, Weizmann Institute of
Science, Rehovot, 7610001, Israel
13PITT PACC, Department of Physics and Astronomy, University of Pittsburgh,
Pittsburgh, PA 15260, USA
14Department of Physics, University of California, Santa Barbara, CA
93106-9530, USA
15Las Cumbres Observatory, 6740 Cortona Dr Ste 102, Goleta, CA 93117-5575, USA
16 Astrophysics Research Centre, School of Mathematics and Physics, Queens
University Belfast, Belfast BT7 1NN, UK
17School of Physics & Astronomy, Cardiff University, Queens Buildings, The
Parade, Cardiff, CF24 3AA, UK
18 SRON Netherlands Institute for Space Research, Sorbonnelaan 2 3584 CA
Utrecht, The Netherlands
19 Department of Astrophysics/IMAPP, Radboud University, P.O. Box 9010, 6500
GL Nijmegen, The Netherlands
20Departamento de Ciencias Fisicas, Universidad Andres Bello, Avda. Republica
252, Santiago, Chile
21Millennium Institute of Astrophysics (MAS), Nuncio Monseñor Sótero Sanz 100,
Providencia, Santiago, Chile
22 The Observatories of the Carnegie Institution for Science, 813 Santa
Barbara Street, Pasadena, CA 91101, USA
E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

Hydrogen-rich, core-collapse supernovae are typically divided into four
classes: IIP, IIL, IIn, and IIb. In general, interaction with circumstellar
material is only considered for Type IIn supernovae. However, recent
hydrodynamic modeling of IIP and IIL supernovae requires circumstellar
material to reproduce their early light curves. In this scenario, IIL
supernovae experience large amounts of mass loss before exploding. We test
this hypothesis on ASASSN-15oz, a Type IIL supernova. With extensive follow-up
in the X-ray, UV, optical, IR, and radio we present our search for signs of
interaction, and the mass-loss history indicated by their detection. We find
evidence of short-lived intense mass-loss just prior to explosion from light
curve modeling, amounting in 1.5 M⊙ of material within 1800 R⊙ of the
progenitor. We also detect the supernova in the radio, indicating mass-loss
rates of $\rm{10^{-6}}$ -10-7 M⊙ yr-1 prior to the extreme mass-loss period.
Our failure to detect the supernova in the X-ray and the lack of narrow
emission lines in the UV, optical, and NIR do not contradict this picture and
place an upper limit on the mass-loss rate outside the extreme period of
$\rm{<10^{-4}}$ M⊙yr-1. This paper highlights the importance gathering
comprehensive data on more Type II supernovae to enable detailed modeling of
the progenitor and supernova which can elucidate their mass-loss histories and
envelope structures and thus inform stellar evolution models.

###### keywords:

supernovae: general supernovae: individual: ASASSN-15oz stars: late-type
stars: winds, outflows techniques: imaging spectroscopy

††pubyear: 2018††pagerange: Signatures of Circumstellar Interaction in the
Type IIL Supernova ASASSN-15oz–B

##  1 Introduction

Stars greater than 8 M⊙ end their lives spectacularly as core-collapse
supernovae (CCSN). CCSNe are divided into two types: Type I, those without
hydrogen in their spectra and Type II, those with hydrogen in their early
spectra (Minkowski, 1941). Although the diversity of Type II light curves was
noted early (Minkowski, 1964; Pskovskii, 1967), it was Barbon et al. (1979)
who first proposed a division of the Type II class into Type IIP SNe (IIP SNe)
and Type IIL SNe (IIL SNe) based on the shape of their light curves. With
small samples, some authors have found a distinction between IIP SNe and IIL
SNe light curves (e.g. Patat et al. 1993, and Patat et al. 1994, Arcavi et al.
2012). However, the recent analysis of large samples of SNe by Anderson et al.
(2014); Faran et al. (2014); Sanders et al. (2015); Valenti et al. (2016);
Galbany et al. (2016), and Rubin & Gal-Yam (2016) show a continuum between the
archetypical IIP and IIL light curves. Given the continuum of light curves
observed, we will use IIP-like to denote SNe at the IIP end of the spectrum
and IIL-like to denote SNe at the IIL end of the spectrum.

From high resolution imaging prior to the SNe explosion, the progenitors of
IIP/IIL SNe are found to be red supergiants (RSG) with masses $\rm{<}$ 17 M⊙
(Smartt et al. 2015 and references there in). Successful modeling of IIP-like
SNe with RSG progenitors has shown that the initial period of decline is due
to adiabatic cooling during the expansion of the ejecta (e.g. Grassberg et al.
1971; Falk & Arnett 1977; Blinnikov & Bartunov 1993; Kasen & Woosley 2009).
When the photosphere reaches $\rm{\sim}$ 6000K, hydrogen recombination
commences, moving the photosphere inward, against the outwardly expanding
ejecta. The rate of the expansion of the ejecta matches the rate of the
recession of the recombination front, which itself is at the constant
temperature of recombination, leading to the almost constant brightness that
define the IIP-like class. When the photosphere passes completely through the
hydrogen dominated layer of the ejecta, the ejecta continues to expand and
cool, as layers rapidly become optically thin until the light curve is powered
by the radioactive decay of
$\rm{{}^{56}Ni\rightarrow{}^{56}Co\rightarrow{}^{56}Fe}$.

It is theorized that IIL-like SN light curves arise from progenitors with
smaller hydrogen envelopes (Grassberg et al., 1971; Young & Branch, 1989;
Blinnikov & Bartunov, 1993). In this scenario, the recombination front would
move faster than the expanding ejecta, causing the brightness to decrease. The
shallow H $\rm{\alpha}$ absorption troughs in IIL-like SNe, possibly due to
less absorbing material, support this picture (Schlegel, 1996; Gutiérrez et
al., 2014). To explode with smaller hydrogen envelopes, IIL-like SN
progenitors must lose mass via stellar winds or mass transfer to a binary
companion. In this paper we will consider the case of mass loss via stellar
winds. As mass-loss increases with zero age main sequence mass (MZAMS) (Heger
et al., 2003; Kasen & Woosley, 2009), this leads to the hypothesis that the
progenitors of IIL-like SNe are more massive than IIP-like progenitors.
Observations that IIL-like SNe have shorter periods before falling to the
radioactive decay phase, are brighter, and have higher velocity ejecta support
this hypothesis (Patat et al., 1994; Gutiérrez et al., 2014; Anderson et al.,
2014; Faran et al., 2014; Sanders et al., 2015; Valenti et al., 2015, 2016).
This idea is intriguing as it solves another outstanding problem, the lack of
progenitors between 17-25 M⊙ (Smartt et al., 2015) (for other possible
solutions see: Davies & Beasor 2018; Walmswell & Eldridge 2012). However,
Valenti et al. (2016) found that there is not a clear difference in the
progenitor population of IIP-like and IIL-like SNe and more importantly, that
the progenitors of IIL-like SNe are not above 17 M⊙. Nevertheless, mass-loss,
whether governed by progenitor mass or some other factor (e.g. metallicity,
rotation, binarity) remains the most promising explanation for the diversity
of light curve shapes.

Mass-loss rates for RSGs are not well understood theoretically and are
extremely difficult to characterize with observations due to both the short
period of this phase as well as the rarity of high mass stars. Thus the mass
loss rates for RSG are uncertain, with typical mass-loss rate estimates
between $\rm{10^{-6}-10^{-4}}$ M⊙$\rm{yr^{-1}}$ and wind velocities between
$\rm{10-100}$ km $\rm{s^{-1}}$ (Mauron & Josselin, 2011). Mass lost from a SN
progenitor over the course of its lifetime will be distributed around the
progenitor with a density that is related to the time the mass-loss occurred,
the mass-loss rate, and the wind velocity. Although interaction has been
observed in SNe at the edge of the distribution in terms of decline rate
(steep) and absolute magnitude (bright; e.g. SN 1979C: Blinnikov & Bartunov
1993; SN 1980K: Chugai 1992; SN 1998S: Chugai 2001; see Section 7), it is
often assumed to be negligible in IIP/IIL SNe. It is possible that the
signatures of interaction with circumstellar material (CSM; with stellar wind
origins) are visible in the light curves and spectra of most IIP/IIL SNe,
especially outside of the optical where there are fewer observations. The
detection and evolution of this interaction can directly probe the mass-loss
history of massive stars and provide valuable information on the nature of
IIP/IIL SN progenitors.

In fact, evidence of interaction, often characterized by a single type of
observation (e.g. X-ray detection), indicates that mass-loss should be
considered in the interpretation of IIP/IIL SNe. Recent light curve modeling
has shown that dense CSM is required to match the early UV and optical light
curves of IIP/IIL SNe with hydrodynamic models of SNe from RSG progenitors
(Gezari et al., 2015; Morozova et al., 2018; Paxton et al., 2018; Förster et
al., 2018). Flash ionization features, the narrow emission lines in Type
IIP/IIL SNe that disappear within the first hours to days post explosion are
also explained by interaction with CSM. These features can be produced either
by episodic high mass-loss rates in single stars (Gal-Yam et al. 2014, e.g.
Yaron et al. 2017) or by lower mass-loss rates in a binary system (Kochanek,
2018). Additional evidence of the importance of mass-loss comes from rare
X-ray and radio observations of IIP/IIL SNe, some of which detect signs of
interaction (e.g. Dwarkadas 2014 and references there in, Morozova & Stone
2018).

Although interaction is seen in some historical IIL-like SNe (e.g.
1979C,1980K, 1998S), as more IIL-like SNe are discovered, it has become clear
that these interacting events are at the edge of the IIL-like distribution
with bright absolute magnitudes and steep decline rates. Furthermore, a large
number of IIL-like SNe, more central to the distribution, lack strong signs of
interaction. Here, we search for interaction in one of these central IIL-like
SN, discovered by the All-Sky Automated Survey for SuperNovae (ASAS-SN),
ASASSN-15oz. With our extensive set of data, ranging from X-ray to radio
wavelengths, we examine different indicators of interaction with CSM to form a
coherent history of mass-loss.

The outline of the paper is as follows. In Section 2 we describe the data
collected at all wavelengths. In Section 3 we present the parameters of
ASASSN-15oz. The photometric and spectroscopic evolution are described in
Sections 4 and 5, respectively. Section 6 describes our search for
interaction, including hydrodynamic light curve modeling and radio analysis.
Finally, we present a comparison to other objects and a unified interpretation
of the mass-loss history indicated by the multi-wavelength observations in
Section 7 and summarize our conclusions in Section 8.

##  2 Observations and Data Reduction

ASASSN-15oz was the closest Type II SN of 2015 and was observed extensively at
all wavelengths. The majority of our observations are provided by the Las
Cumbres Observatory (LCO; Brown et al. 2013) Supernova Key Project (2014-2017)
and the Public ESO Spectroscopic Survey for Transient Objects (PESSTO). We
supplemented these with observations in the X-ray, ultraviolet (UV), optical,
near infra-red (NIR), and radio. The photometric and spectroscopic data
reduction is described in this section. A table of spectroscopic observations
is presented in Table 4. A complete list of photometric observations is given
in the Supplemental Material. Table 5 gives an example of the photometric
observation table format. All photometric data can be obtained from the
SNDAVIS database111http://dark.physics.ucdavis.edu/sndavis/transient and all
spectroscopic data can be found on WISeREP222https://wiserep.weizmann.ac.il.

###  2.1 Optical

The optical light curve was closely monitored by the LCO from discovery
through day 400 in the filters B, V, g, r, and $i$. Unfortunately, the SN
passed behind the sun around day 80, leaving a gap in the observations from
day 87 to day 179, most notably missing the fall from plateau. Figure 1 shows
the complete multi-band light curve. LCO observations were reduced with the
LCO imaging pipeline, lcogtsnpipe (Valenti et al., 2016). This pipeline
employs PSF photometry, removing background contamination by fitting a low
order polynomial to the host galaxy. Instrumental magnitudes were converted to
apparent magnitudes using stars in the the
APASS333https://www.aavso.org/apass, Sloan, and Landolt catalogs. Apparent
magnitudes for the U, B, V, R, and I filters are given in the Vega magnitude
system and apparent magnitudes for the g, r, and i filters are given in the AB
magnitude system.

Figure 1: The complete UV, optical, and NIR light curve of ASASSN-15oz
expressed in terms of days since explosion. Apparent magnitude (+offset) is
plotted on the left and absolute magnitude (+offset) is plotted on the right.
Each color corresponds to a single filter, labeled in the legend. The optical
Swift filters are denoted with an "s" in the name (e.g. us) and plotted with
the triangle symbol. NIR observations are shown with the diamond symbol. To
differentiate the U, B, V, R, and I filters and the g, r, and i filters, we
use a square symbol for the former and a circle for the later. Each filter has
been shifted by a global offset (denoted in the legend) for viewing purposes.

Extensive optical spectroscopy was obtained during the photospheric phase with
the FLOYDS spectrograph (Brown et al., 2013) on the 2m LCO telescope at Siding
Springs, Australia and the EFOSC2 spectrograph (Buzzoni et al., 1984) on the
3.6 m NTT telescope at La Silla, Chile during the photospheric and nebular
phase. One-dimensional spectra were extracted and calibrated using the FLOYDS
pipeline (Valenti et al., 2014) and the EFOSC tasks in the PESSTO pipeline
(Smartt et al., 2015). Both of these pipelines combine IRAF tasks to bias
subtract and flat field the data and locate, extract, wavelength calibrate,
and flux calibrate the one-dimensional spectrum.

Data were also taken using the X-Shooter echelle spectrograph (Vernet et al.,
2011) on the 8.2m Very Large Telescope (VLT) at the European Southern
Observatory (ESO). X-Shooter has three arms (UVB, VIS, NIR) which combined
provide continuous wavelength coverage over $\rm{3000-24,800}$ Å. For our
data, we used slit widths of 1.0, 0.9, and 0.9 arcsec in the UVB, VIS, and NIR
arms, corresponding to resolutions of R $\rm{\sim}$5400, 8900, and 5600,
respectively. The UVB and VIS data were reduced via a modified version of the
EsoReflex pipeline (Freudling et al., 2013), with improvements to the sky-
subtraction and rebinning procedure on the highly dispersed echelle spectrum.

Figure 2 shows the evolution of the optical spectra over time. The progression
starts with a blue spectrum with broad hydrogen lines at the top of the
figure. Over time (moving down the figure), the spectra develop metal features
such as iron and scandium as the velocities decrease and the line widths
narrow.

We obtained late time spectra from the Gemini South Observatory through
program GS-2016A-Q-75-25 (PI: Valenti). The Gemini spectra of ASASSN-15oz were
taken on day 287 (2016 Jun 09), day 288 (2016 Jun 10), and day 290 (2016
Jun12) using the Gemini Multi-Object Spectrograph (GMOS) (Hook et al., 2004).
Data were taken with a blue setting with the B600 grating, no filter, a
1.0″slit, and 2x2 binning, and a red setting with the R400 grating, OG515
order blocking filter, a 1.0″slit, and 2x2 binning. These settings provide a
wavelength range 3630 - 6830 Å with a dispersion of 1.0 Å $\rm{pixel^{-1}}$ in
the blue and a wavelength range 5030 - 9800 Å with a dispersion of 1.3Å
$\rm{pixel^{-1}}$ in the red. We used the Gemini iraf package for GMOS to
process our data. The raw two-dimensional images were bias subtracted,
overscan corrected, flat-fielded, wavelength calibrated using CuAr emission
lamps, cleaned of cosmic rays, and background subtracted prior to extraction
of the 1D target spectra. The spectra were then calibrated using a sensitivity
function derived from a standard star observed for the program.

Figure 2: A time series of the photospheric spectra of ASASSN-15oz at rest
wavelengths. The phase of each spectrum is marked on the right. Potential
telluric contamination is marked with crossed circles above the day 24.5
spectrum. The first spectrum taken near maximum light shows hydrogen features
(identified at the top of the figure). At this early phase, the cachito
feature (see Section 6.2) is already visible. Over time the cachito feature
fades while the hydrogen emission grows and metal lines become visible and
grow in strength. These lines are identified in the bottom spectrum.

###  2.2 UV and X-Ray Observations

UV and X-Ray imaging observations were obtained concurrently with Swift from
day nine (2015 Sep 05) through day 83 (2015 Nov 18). In addition to the
imaging observations, two epochs were observed with the Swift grism under
program 1114241 (PI: Valenti). Imaging observations were processed with
aperture photometry following Brown et al. (2009) and employing the updated
zeropoints of Breeveld et al. (2010). UV grism observations were taken at two
roll angles and extracted using the UVOTPY pipeline444Kuin (2014) (Kuin et
al., 2015). There is a nearby star coincident with the SN at one roll angle
and just above it at the other roll angle. Although we tried custom
extractions using different extraction box height and template subtraction
using TRUVOT (Smitka, 2015) we were unable to produce a better calibration
than the default pipeline. Due to the distance of the ASASSN-15oz, only the
first of the two spectroscopic epochs has sufficient S/N to be extracted.
These spectra are shown in Figure 3.

Figure 3: The Swift UV spectra from day nine (2015 Sep 05), boxcar smoothed
with a kernel of three pixels. The spectrum taken with a roll angle of PA =
248∘ is in pink, the spectrum taken with a roll angle of PA = 260∘ in indigo,
and average flux is in black. We tentatively identify UV FeII, TiII, and MgII
features as well as mark the Balmer lines. These spectra are not template
subtracted, therefore they show residuals from zeroth order light. These
features can be identified as being in the spectrum of one roll angle but not
the other. The lack of narrow line emission indicates a lack of strong
interaction.

Aperture photometry was performed on the X-ray observations using the HEAsoft
packages xselect (Blackburn, 1995) and xspec; (Arnaud, 1996). The background
was selected as a region with no x-ray sources in the image. The source was
extracted using an 18″aperture. Count rates were converted to a luminosity
using PIMMS (Mukai, 1993) assuming a power-law model with a photon index of
two. Integrating over all available SWIFT/XRT epochs, no significant detection
was made after background subtraction.

###  2.3 NIR Observations

ASASSN-15oz was observed in the NIR with the SOFI instrument on the NTT
telescope through the PESSTO collaboration, the SpeX instrument (Rayner et
al., 2003) on the 3.0-m NASA Infrared Telescope Facility (IRTF), and the NIR
arm of the X-Shooter spectrograph. The data reduction of the NIR arm of the
X-Shooter spectrum is described with the UV and VIS arms in Section 2.1.

The SOFI spectroscopic data were reduced with the PESSTO package (Smartt et
al., 2015). Photometric measurements were performed with the QUBA pipeline
(Valenti et al., 2011), which performs DAOPHOT-based (Stetson, 1987) point-
spread-function fitting photometry on the SN and on the selected reference
stars. J, H, and K photometry was then calibrated to the 2MASS magnitude
system.

The SpeX data were taken in so-called SXD mode, where the spectrum is cross-
dispersed to obtain wavelength coverage from $\rm{\sim}$0.8–2.4 $\rm{\mu}$m in
a single exposure, spread over six orders. All observations were taken with
the slit aligned along the parallactic angle, and we employed a classic ABBA
technique for improved sky subtraction. HD 177074, an A0V star, was observed
adjacent in time to the science observations for flux and telluric
calibration. The spectrum was reduced in a standard way using the publicly
available Spextool software package (Cushing et al., 2004) and corrections for
telluric absorption utilized XTELLCOR (Vacca et al., 2003) and the A0V star
observations.

The three photometric observations from day eight (2015 Sep 04), day 38 (2015
Oct 04), and day 86 (2015 Nov 21) are shown in Figure 1. Spectroscopic
observations from SOFI (day nine; 2015 Sep 05 and day 39; 2015 Oct 05),
X-Shooter (day 26; 2015 Sep 21), and SpeX (day 44; 2015 Oct 10) are shown in
Figure 4.

Figure 4: The evolution of the infrared spectra at rest wavelengths during the
photospheric phase. The spectral evolution in the IR mirrors the optical
evolution, with broad hydrogen features narrowing and metal lines developing
over time. The first spectrum (black) taken eight days post-explosion shows
broad hydrogen features. The later spectra at day 38 (indigo) and 43 (yellow)
reveal the development metal lines. Although contaminated with narrow telluric
features, the high resolution X-Shooter spectrum (pink) shows an intermediate
phase. Regions of severe sky contamination are marked with grey boxes and are
masked out of the specta.

###  2.4 Radio Observations

We observed ASASSN-15oz using the NSF’s Karl G. Jansky Very Large Array (VLA)
(VLA/15B-362, PI: Valenti) on day 28 (2015 Sep 24) and day 42 (2015 Oct 08).
In both of the observations we used J1924-2914 and 3C48 as a phase calibrator
and a flux calibrator, respectively. We used the Common Astronomy Software
Applications (CASA) standard packages and pipelines to perform the data
calibration and imaging. The SN was detected in both observation with the
following flux measurements: September 24: $\rm{120\pm 23\,\mu}$Jy at
$\rm{4.8}$ GHz (C-band) and a $\rm{3\,\sigma}$ non-detection limit at
$\rm{22}$ GHz (K-band); October 08: $\rm{210\pm 21\,\mu}$Jy at $\rm{4.8}$ GHz
(C-band) and a $\rm{80\pm 17\,\mu}$Jy at $\rm{15}$ GHz (KU-band).

##  3 Supernova Parameters

ASASSN-15oz was discovered by the ASAS-SN team on 2015 Aug 31.09 (JD
2457265.59; V = 14.6 mag; RA = 19:19:33.55, Dec=-33:46:01.2) and announced on
2015 Sep 03 (JD 2457268; Brown et al. 2015). The next day, 2015 Sep 04 (JD
2457270.0), LCO classified it as a Type II SN one week after explosion
(Hosseinzadeh et al., 2015). Figure 5 shows the Digital Sky Survey (DSS) image
of the low luminosity host galaxy, HIPASS J1919-33, prior to explosion (top
panel) and our first V-band image of ASASSN-15oz taken on day eight (2015 Sep
04). Imaging of the host galaxy, HIPASS J1919-33, from 2015 Aug 23 (JD
2457257.56) does not show any evidence of ASASSN-15oz (V $\rm{>}$ 17.8; Brown
et al. 2015). We take this date to be a lower limit on the explosion epoch and
define the mid-point between the non-detection and the detection as explosion
date: $\rm{t_{expl}=}$ 2015 Aug 27 $\rm{\pm 4}$ days (JD 2457261.5). This is
in agreement with the value used in Gutiérrez et al. (2018).
When non-detection observations are not available, other methods can be used
to determine the explosion epoch. Gutiérrez et al. (2017) found that the
explosion epoch can be determined, with an average errors of 4.9 days, by
fitting the blue part of the first spectrum with SNID555Blondin & Tonry (2011)
(Blondin & Tonry, 2007) and using the explosion epoch of the best fit
template. This method relies on the existence of a template SN that is similar
in both phase and type to the object being fit. ASASSN-15oz is a IIL-like SN.
SNID has very few templates for IIL-like SNe, even after adding the new
templates provided by Gutiérrez et al. (2017). Therefore, in addition to the
Gutiérrez et al. (2017) templates, we add templates for two well observed
supernovae with well constrained explosion epochs: SN 2012A (Tomasella et al.,
2013), and SN 2013ej (Childress et al., 2016; Dhungana et al., 2016; Valenti
et al., 2014; Smartt et al., 2015). With these templates, we confirm our
explosion epoch by fitting the first spectrum (2015 Sep 04) using SNID v. 5.0.
The five best fit spectra imply an explosion epoch of 2015 Aug 22$\rm{\pm}$4.9
(JD 2457256.5), consistent with the range we propose based on non-detections.
Interestingly, the hydrodynamic light curve modeling also places the best fit
explosion epoch earlier than 2016 Aug 27 (see Section 4). As the uncertainties
of both these methods are greater than that of the non-detection, throughout
this paper we adopt the value determined by the non-detection as the explosion
epoch, $\rm{t_{expl}=}$ 2015 Aug 27 $\rm{\pm 4}$ days (JD 2457261.5), unless
otherwise specified.

We use the Hubble flow distance modulus to the host galaxy, HIPASS J1919-33,
$\rm{\mu=32.3\pm 0.2}$ mag, with H 0=73 km $\rm{s^{-1}}$ $\rm{Mpc^{-1}}$
corrected for Virgo infall 666The NASA/IPAC Extragalactic Database (NED) is
operated by the Jet Propulsion Laboratory, California Institute of Technology,
under contract with the National Aeronautics and Space Administration..

Figure 5: Top: The host galaxy of ASASSN-15oz, HIPASS J1919-33, taken prior to
the SN explosion by DSS. Bottom: V-band image of ASASSN-15oz taken by LCO
eight days post explosion (2015 Sep 04). The location of the ASASSN-15oz is
marked by the pink ticks in both images.

We use the VIS arm of the X-shooter spectrum (R $\rm{\sim}$8900) to constrain
the host galaxy reddening and to confirm the galactic reddening using the
equivalent widths of the Na I D1 and D2 lines ($\rm{\lambda}$5890,
$\rm{\lambda}$5896) (Poznanski et al., 2012). We find no evidence of Na I
absorption at the host redshift and therefore assume no host reddening (see
Figure 6). By fitting gaussian profiles to the galactic Na I absorption we
find $\rm{E(B-V)=0.12\pm 0.07}$ mag, consistent with the value given by
Schlafly & Finkbeiner (2011). We adopt the Schlafly & Finkbeiner (2011) value
of $\rm{E(B-V)=0.08}$ mag for the analysis in this paper and use the
extinction law of Cardelli et al. (1989) with $\rm{R_{V}=3.1}$.

Figure 6: The galactic (solid gray lines) and host (dashed gray lines) Na I D1
and D2 lines in the high resolution X-shooter spectrum from day 26; (2015 Sep
21). The galactic profiles are consistent with the galactic extinction of
Schlafly & Finkbeiner (2011) while the absence of absorption at the host
redshift indicates a lack of dust in the host galaxy.

##  4 Light Curve Evolution

The light curve of ASASSN-15oz reaches peak brightness in the V-band 8.25 days
after explosion, one day after the start of our observations, reaching an
above average absolute magnitude of -18.05 $\rm{\pm}$ 0.025 mag. The early
light curve rises until around day 20 in the redder filters, flattening in the
V-band and falling in the blue bands during this same period. It then falls in
all bands for the remainder of the observations, changing slope around day 40
and again sometime after day 85, as it enters the radioactive decay phase,
following the fall from plateau. Following Anderson et al. (2014), we use the
V-band light curve to fit a slope to the first steep initial decline after
maximum ($\rm{s_{1}}$) and to the second shallower slope prior to the fall
from plateau ($\rm{s_{2}}$). We also fit a global decline rate, $\rm{s_{50v}}$
from soon after maximum to prior to the fall from plateau. During the nebular
phase we fit a slope to the radioactive decay tail ($\rm{s_{tail}}$777this is
equivalent to $\rm{s_{3}}$ in Anderson et al. (2014)). The slopes we fit as
well as the details of the fit are summarized in Table 1 and plotted on top of
the V-band light curve in Figure 7. Both the high slopes and bright maximum
luminosity of ASASSN15oz are similar to those of other IIL-like SNe.

Table 1: The best-fit slope to the V-band light curve of ASASSN-15oz measured between the start and end phase listed in the table. All slopes are measured in units of magnitudes per 50 days. Slope Type | Slope (mag $\rm{(50days)^{-1}}$) | Slope Error (mag $\rm{(50days)^{-1}}$) | Start Phase (day) | End Phase (day)
---|---|---|---|---
$s_{1}$ | 1.21 | 0.066 | 21.0 | 38.0
$s_{2}$ | 0.99 | 0.021 | 45.0 | 78.0
$s_{50V}$ | 1.09 | 0.014 | 18.4 | 78.0
$s_{tail}$ | 0.70 | 0.029 | 206.9 | 368.0

The slope of the radioactive decay tail is steeper than expected from the
radioactive decay of $\rm{{}^{56}Ni\rightarrow{}^{56}Co\rightarrow{}^{56}Fe}$
implying that the ejecta is not fully trapping and reprocessing the gamma rays
produced by the radioactive decay.

Figure 7: The $\rm{s_{1}}$ (pink), $\rm{s_{2}}$ (yellow), $\rm{s_{50V}}$
(indigo), and $\rm{s_{tail}}$ (black) fits plotted over the period used to
derive the relation. The V-band light curve is plotted in gray. Although
errors in the apparent magnitude are plotted, in many cases these are smaller
than the points.

The pseudo-bolometric luminosity is computed following Valenti et al. (2008).
To summarize the method, the apparent magnitudes are corrected for galactic
extinction. We use the U, B, g, V, r, R, i, and I-bands to calculate the
psuedo-bolometric luminosity at the V-band cadence. The flux density is
calculated from the apparent magnitude at the effective wavelength of each
filter and integrated using Simpson’s Rule. Finally, the flux is converted to
a luminosity using the distance modulus. When V-band observations extended
either earlier or later than a filter, we use the color with the neighboring
filter to extrapolate the flux.

###  4.1 Nickel mass determination

The amount of $\rm{{}^{56}Ni}$ synthesized during the explosion can be
inferred using the late time ($\rm{>}$ 150 days) photometry. At this time, the
energy output is powered by the radioactive decay of
$\rm{{}^{56}Ni\rightarrow{}^{56}Co\rightarrow{}^{56}Fe}$ which produces gamma
rays that are reprocessed by the ejecta producing optical light. Assuming that
the SN has the same SED as the extremely well characterized SN 1987A, the
$\rm{{}^{56}Ni}$ mass synthesized in the SN explosion can be found by scaling
the pseudo-bolometric luminosity of SN when it is powered by $\rm{{}^{56}Co}$
to that of SN 1987A with the following equation (Spiro et al., 2014):

$M({}^{56}Ni)=0.075M_{\odot}\times\frac{L_{SN}(t)}{L_{87A}(t)}$ (1)

where $\rm{M(\,{}^{56}Ni)}$ is the synthesized $\rm{{}^{56}Ni}$ mass,
$\rm{L_{SN}(t)}$ is the pseudo-bolometric luminosity of the SN at time t, and
$\rm{L_{87A}(t)}$ is the pseudo-bolometric luminosity of SN 1987A at time t
using the same filters as were used to compute the pseudo-bolometric
luminosity of the SN.

This scaling can be performed at any time, as long as the slope is the same as
that of SN1987A, i.e. consistent with the complete trapping of gamma rays. As
the SN ages, the optical depth of the gamma rays decreases. Thus a SN is most
likely to have complete trapping immediately after the fall from plateau.
Unfortunately, for ASASSN-15oz, there is not complete trapping for any part of
the observed radioactive decay tail and there are no observations of the
period immediately after the fall from plateau. We therefore consider the
maximum and minimum $\rm{{}^{56}Ni}$ mass allowed by the observations, under
the assumption that there is complete trapping immediately following the fall
from plateau. For both cases we also assume that the fall from plateau takes
20 days, based on light curve fits presented in Valenti et al. (2016). An
upper limit on the $\rm{{}^{56}Ni}$ mass is obtained by assuming ASASSN-15oz
fell from plateau immediately after the last photospheric point (see upper
panel of Figure 8). Scaling the psuedo-bolometric luminosity of SN1987A at 102
days to that of ASASSN-15oz at the same time, we calculate a $\rm{{}^{56}Ni}$
mass of 0.11 M⊙. To find the lower limit, we estimate a conservative value of
the time to the middle of the fall from plateau ($\rm{t_{pt}}$) to be 125 days
using Figure 5a in Valenti et al. (2016). Using this $\rm{t_{pt}}$ we
calculate the fall from plateau to end 135 days after explosion and
extrapolate the tail fit to this phase (see lower panel of Figure 8). We
compute a $\rm{{}^{56}Ni}$ mass of 0.08 M⊙ by scaling the psuedo-bolometric
luminosities at 135 days. In both panels of Figure 8, the observations are
represented with black circles, the $\rm{s_{2}}$ slope with a pink line, the
fall from plateau with a indigo line, the tail slope with a yellow line, and
the expected tail slope from complete trapping is marked with a dashed green
line. The time that is used to scale SN 1987A to ASASSN-15oz is marked with a
cyan star.

Figure 8: The artificial light curve of ASASSN-15oz for the shortest (top) and
longest (bottom) plateau lengths. The black points plot the bolometric
luminosity while the pink line represents a fit to the s2 data, the indigo
line is the artificial fall from plateau, and the yellow line is a fit to the
radioactive decay tail. The cyan star marks the point to which the light curve
of SN 1987A is scaled to derive a $\rm{{}^{56}Ni}$ mass. Plotted as a green
dashed line is the slope of the radioactive decay tail if there was complete
trapping. An upper limit on the $\rm{{}^{56}Ni}$ mass of 0.11 M⊙ is derived
from the light curve in the top panel while a lower limit of 0.08 M⊙ is
derived from the lower panel.

##  5 Spectroscopic Analysis

The spectroscopic evolution of a SN during the photospheric phase provides
insight into details of the SN ejecta. The changes in velocity and the
individual line profiles describe the geometry and energetics of the ejecta.
The presence of different species at different times gives information about
the chemical composition, temperature, and density of the ejecta. In this
section we analyze optical and NIR spectra.

###  5.1 Optical Evolution

Initial line identification was performed by comparing the spectra to that of
SN 1999em (Leonard et al., 2002). Identification was confirmed using the
spectrum synthesis code SYN++ 888Thomas (2013)(Thomas et al. 2011; parameter
details in Table 6). The results of the SYN++ fit can be seen in Figure 9.

Figure 9: A comparison of the model spectra produced by Syn++ (colored lines)
and the observed spectrum (black) from day 40 (2015 Oct 06). The top panel
shows the best fit with all elements while the remaining panels show the fit
for each element individually. We fit H $\rm{\alpha}$ and H $\rm{\beta}$
separately as H $\rm{\alpha}$ is not well modeled by the pure resonance
scattering assumed in Syn++. The H $\rm{\beta}$ fit is used in the combined
spectrum in the top panel.

We use two methods to find the velocity of lines for which individual
components can be resolved, one for blended features, and one for individual
features. For lines that are not blended or that are dominated by the
absorption of one ion (e.g. Sc II ($\rm{\lambda}$ 5526, 5662, 6262)), we
employ the method presented in Silverman et al. (2012). For blended features
we fit a multi-component gaussian to the feature. We use this method to fit H
$\rm{\alpha}$ ($\rm{\lambda}$ 6561), O I ($\rm{\lambda}$ 7774), FeII
($\rm{\lambda}$ 5169), and Ca II ($\rm{\lambda\lambda\lambda}$ 8498, 8542,
8662). Details of both fitting procedures are given in Appendix B. Rather than
expressing an uncertainty on the fit, we calculate the range of velocities
that encompass 68.2% of the integrated (continuum subtracted) flux around the
feature minimum. Because each line is formed in an extended region of the
ejecta, this represents the range of velocities at which each line forms.

Figure 10 shows the measured velocities. The distribution of velocities is
typical of other IIP/IIL SNe representing the distribution of elements
throughout the ejecta. A comparison of the velocities derived in this paper to
the average values of 122 IIP/IIL SNe described in Gutiérrez et al. (2017) is
shown in Figure 11. The shaded regions represent the standard deviation of the
sample. We find above average velocities throughout the ejecta.

Figure 10: The evolution of the velocity of H $\rm{\alpha}$, H $\rm{\beta}$,
Na I, O I, Ca II, Sc II, and Fe II over time. The points represents the
minimum of the feature showing the velocity with the largest optical depth
while the bars plot the range of velocities found in the line forming region.

The coincidence of velocities of different species, especially at late times,
indicates that mixing may have been important for ASASSN-15oz.

Figure 11: A comparison of the velocity of ASASSN-15oz (pink) to the mean
velocity of 122 IIP/IIL SNe (Gutiérrez et al., 2017) (black) for H
$\rm{\alpha}$ (left), H $\rm{\beta}$ (middle), and Fe II ($\rm{\lambda 5169}$;
right). The gray regions represent the standard deviations of the mean
velocities. The velocity of ASASSN-15oz is $\rm{>1\sigma}$ above average for
all three ions. Each of these line originates in a different part of the
ejecta indicating that this is a global trend and the explosion energy is
above average. Following (Silverman et al., 2012) we select velocity errors of
2 Å. These are contained within the symbols and so are not plotted.

###  5.2 Infrared Evolution

Line identification was performed by comparing the NIR SOFI and SpeX spectra
to those in Valenti et al. (2015) and Tomasella et al. (2018). Figure 4 shows
the NIR spectra with prominent features labeled. Like the optical, the first
spectrum, taken eight days post explosion (2018 Sep 05), is dominated by
hydrogen features. The later spectra also mirror the optical evolution,
developing metal lines as the ejecta slows and cools. We use the IR spectra to
search for evidence of interaction via a high velocity He I feature (see
Section 6.2).

###  5.3 Progenitor Mass from Nebular Spectra

Around 200 days the SN ejecta becomes optically transparent, revealing the
inner core. At this time the SN enters the nebular phase, showing strong
emission lines and no continuum (see Figure 12). Analysis of the spectra at
this time can constrain the geometry of the explosion as well as the abundance
of different elements. In particular, spectral modeling of this stage is a
powerful tool to constrain the nature of the progenitor. Jerkstrand et al.
(2012); Jerkstrand et al. (2014) have shown that the intensities of a few
lines are sensitive to the MZAMS of the progenitor. Specifically, there is a
tight monotonic correlation between the strength of the [OI]
($\rm{\lambda\lambda 6300,6334}$) line and the MZAMS of the progenitor. This
is because [OI] is more isolated and less sensitive to the explosive
nucleosynthesis of the SN than the other lines and thus reflects the oxygen
abundance of the progenitor which has been shown to correlate well with
progenitor mass (Woosley & Weaver, 1995).

Jerkstrand et al. (2014) modeled nebular spectra for 12, 15, 19, and 25 M⊙
progenitors ($\rm{M_{ZAMS}}$). Starting with the explosion models from Woosley
& Heger (2007), the ejecta is divided into zones based on chemical composition
and the evolution of the spectrum is found by modeling the gamma ray transport
and deposition, non-thermal electron degradation, NLTE ionization and
excitation, and Monte Carlo radiative transfer.

Four nebular spectra of ASASSN-15oz were taken between day 228 (2016 Apr 11)
and day 389 (2016-09-19). The first three spectra are high S/N and are use for
modeling. In Figure 12 we compare these epochs to the different mass models,
scaling the models to the observations over the full wavelength range. A zoom
in on the [OI] line is displayed in the inset of each panel. We find the [OI]
strength of ASASSN-15oz falls between the 15 and 19 M⊙ models, consistent with
the 17 M⊙ found by modeling the light curve (see Section 6.4).

As a sanity check of the comparison between observations and synthetic
spectra, we infer the $\rm{{}^{56}Ni}$ mass from the scale factor we used to
scale the models to the observations. Following Jerkstrand et al. (2018), the
scaling combined with the equation for the luminosity of the cobalt decay
(Equation 6 of Jerkstrand et al. 2012) yields the $\rm{{}^{56}Ni}$ mass
(assuming complete trapping):

$\frac{F_{obs}}{F_{mod}}=\frac{d_{mod}^{2}}{d_{obs}^{2}}\frac{(M_{56,Ni}){}_{obs}}{(M_{56,Ni}){}_{mod}}exp\left(\frac{t_{mod}-t_{obs}}{111.4}\right)$
(2)

where $\rm{F_{obs}}$ and $\rm{F_{mod}}$ are the observed and model fluxes,
$\rm{d_{obs}}$ and $\rm{d_{mod}}$ the observed and model distances,
$\rm{(M_{56,Ni}){}_{obs}}$ and $\rm{(M_{56,Ni}){}_{mod}}$ are the observed and
model $\rm{{}^{56}Ni}$ masses, and $\rm{t_{obs}}$ and $\rm{t_{mod}}$ are
phases of the observation and the model. Using this relation, we average the
$\rm{{}^{56}Ni}$ mass found for each mass model to find an $\rm{{}^{56}Ni}$
mass of 0.03 M⊙ for day 228 (2016 Apr 11) and day 288 (2016 Jun 10) and 0.02
M⊙ for day 340 (2016 Aug 03). These values are a factor of four smaller than
our lower limit for the $\rm{{}^{56}Ni}$ mass (see Section 4.1). This is due
at least in part to the fact that the spectral synthesis modeling does not
account for the incomplete gamma ray trapping in ASASSN-15oz. The fact that
several IIP/IIL SNe do not have complete trapping (see Section 7) should be
taken into account in the future.

Figure 12: The nebular spectra of ASASSN-15oz from day 228 (2016 Apr 11; top),
day 288 (2016 Jun 10; middle), and day 342 (2016 Aug 03; day 340; bottom)
plotted with the scaled models of Jerkstrand et al. (2014) for $\rm{M_{ZAMS}}$
= 12, 15, 19, and 25 M⊙. Most features in the observed spectra are well
matched by the models. The inset shows the [O I] ($\rm{\lambda\lambda
6300,6334}$) doublet, the strength of which increases monotonically with
progenitor mass. The [O I] flux falls between the 15 and 19 M⊙ models,
consistent with the 18 M⊙ progenitor mass found using light curve modeling in
Section 6.4.

##  6 Evidence for interaction

Although it is well known that massive stars undergo significant mass-loss
during their lives, the details of this mass-loss (timing, rate, speed,
geometry) are not well understood. The location and density of circumstellar
material around a star record the amount of mass lost at different times
during a star’s evolution. As the SN shock passes through this material, it
provides a unique probe of this material and thus the mass-loss of the
progenitor. While the mass-loss of stars with dense CSM (e.g. IIn SNe) that
show strong evidence of interaction, such as narrow emission lines in
photospheric spectra or irregular photometric brightening, is often studied,
the mass-loss of SNe with lower density CSM is often ignored.

In IIP-like SNe, the plateau is the result of the recession of the photosphere
due to hydrogen recombination matching the expansion of the SN ejecta. The
linear decline of the light curves of IIL-like SNe is then explained by the
photosphere receding faster than the ejecta’s expansion. The photosphere’s
recession is governed by the rate at which the ejecta becomes optically thin.
This rate is determined by either the expansion velocity of the ejecta or the
amount of hydrogen in the ejecta, or both. If the expansion velocity is above
average, the ejecta becomes optically thin faster, leading to a more IIL-like
light curve. Likewise, if there is less hydrogen in the ejecta, then the
recombination front moves more rapidly through the ejecta, leading to a more
IIL-like light curve. IIL-like SNe show above average ejecta velocities and
weaker than average hydrogen absorption (Gutiérrez et al., 2014), implying
that both of these effects are at play. Therefore, the linear shapes of the
light curves of IIL-like SNe indicate that their progenitors have smaller
hydrogen envelopes. This then implies that IIL-like SNe should undergo more
mass-loss than IIP-like SNe. ASASSN-15oz, a IIL-like SN with photometric and
spectroscopic observations from the X-ray through the radio is the ideal SN on
which to study the importance of mass-loss for IIP-like and IIL-like SNe. The
different epochs and wavelengths probe different periods of mass-loss allowing
for the characterization of the mass-loss history of the progenitor. In this
section we analyze the evidence for interaction at different wavelengths. We
note that we cannot search for flash ionization features as there are no
spectroscopic observations until near maximum.

###  6.1 UV and X-Ray

ASASSN-15oz is not detected in any of the Swift X-ray observations. To
determine an upper limit on the X-ray luminosity, we sum all observations and
find a limiting count rate of $\rm{5.775\times 10^{-4}}$ counts $\rm{s^{-1]}}$
(0.4-5 keV). Assuming Galactic absorption of 6.18x$\rm{10^{20}}$ cm-2
(Kalberla et al., 2005), this corresponds to an unabsorbed flux of
$\rm{1.647\times 10^{-14}}$ erg cm-2 $\rm{s^{-1}}$ and adopting a distance of
28.83 Mpc, a luminosity upper limit of $\rm{2.56\times 10^{39}}$erg
$\rm{s^{-1}}$.

We compare this upper limit to the model of SN 1999em (Chugai et al., 2007)
which has similar progenitor and explosion parameters. They model the
interaction as an infinitely thin double-shock structure (Chevalier, 1982;
Nadyozhin, 1985) composed of a forward shock propagating into the surrounding
gas and reverse shock propagating backwards into the SN ejecta, towards its
core. Chugai et al. (2007) find that the reverse shock dominates the X-ray
observations for a pre-explosion mass loss rate of $\rm{10^{-6}}$
M⊙$\rm{yr^{-1}}$ and a wind velocity of $\rm{10}$ km $\rm{s^{-1}}$ producing
luminosities between $\rm{10^{38}}$ and $\rm{10^{39}}$ erg $\rm{s^{-1}}$. The
forward shock produces luminosities comparable to the reverse shock in the
first 10 days and decreasing to below $\rm{10^{34}}$ erg $\rm{s^{-1}}$ by day
100. Our non-detection is at the upper limit of this range, making it
plausible that if there is X-ray emission, it is below our detection limits.

We use SYN++ to model and tentatively identify absorption features from Fe II,
Ti II, and Mg II features in the first Swift UV grism spectrum (see Figure 3).
This spectrum shows no sign of the narrow emission lines typically associated
with strong interaction ( $\rm{\dot{M}>10^{-4}}$ g $\rm{cm^{-1}}$; Kiewe et
al. 2012)

###  6.2 Optical Spectroscopy

Chugai et al. (2007) proposed a high velocity hydrogen and helium line are
produced by interaction of the SN ejecta with the stellar wind. According to
their model, a notch on the blue side of the H $\rm{\alpha}$ profile should
appear 40-80 days post explosion. Although they could construct a model with
the introduction of a cold dense shell in which a similar notch was visible in
H $\rm{\beta}$ their simplest model did not show H $\rm{\beta}$ due to its low
opacity in the wind. They do however, predict absorption in He I $\rm{\lambda
10830}$ 20-60 days post explosion, a line that is only excited when wind is
present.

A feature blueward of H $\rm{\alpha}$ is often present in Type II SNe,
however, it is unclear if this feature is due to high velocity hydrogen or to
Si II ($\rm{\lambda}$6355) and it has thus been named the cachito feature to
avoid association with a particular species (Gutiérrez et al., 2017). Building
on a the analysis of smaller samples (e.g. Leonard et al. 2002, Inserra et al.
(2013), Valenti et al. 2014), Gutiérrez et al. (2017) searched for this
feature in 122 IIP/IIL SNe. They find the cachito feature in 70 SNe in their
sample and divide the detections in early phase detections (phase < 40 days)
and late phase detections (phase > 40 days). The velocities of the cachito
feature observed in the early sample are most often well matched to the Fe II
velocity if interpreted as Si II. For this reason, they interpret the cachito
feature in the early sample as Si II and at late times as high velocity
hydrogen. Figure 13 shows the time evolution of the region surrounding H
$\rm{\alpha}$ (left panel), H $\rm{\beta}$ (middle panel), and He I (right
panel). H $\rm{\alpha}$ and H $\rm{\beta}$ are marked with a brick dotted line
in the left and middle panels, respectively. The dashed line in the left panel
marks the location of the cachito feature and in the middle panel, the
expected location of the cachito feature in H $\rm{\beta}$. The time of
expected CSM interaction is marked in grey. ASASSN-15oz shows a strong cachito
feature in the first spectrum, eight days post explosion. This feature
strengthens until about day 15 and then fades until it is no longer visible at
day 60. We find no evidence for a high velocity feature in H $\rm{\beta}$. The
He I feature is fully blended with C I ($\rm{\lambda 10691}$) and
Paschen-$\rm{\gamma}$ and if it is present, it is impossible to deblend and
identify.

Figure 13: The evolution of the cachito feature (left panel; dashed line) over
time. The feature is visible bluewards of H $\rm{\alpha}$ (left panel; dotted
line) in the first spectrum, eight days post explosion. It further increase in
strength for the next 10 days, then decreases until it is barely visible at 60
days post explosion. This evolution is counter to the evolution described in
Chugai et al. (2007) in which the high velocity hydrogen feature becomes
visible around day 40, making it unlikely that it is high velocity hydrogen.
There is no evidence of a high velocity hydrogen feature in H $\rm{\beta}$
(center panel; H $\rm{\beta}$ marked with dotted line; cachito velocity marked
with dashed line) although this is not surprising given the low opacities
predicted by Chugai et al. (2007). The high velocity He I
($\rm{\lambda=10830}$) feature (right panel) also predicted by the models of
Chugai et al. (2007) is too heavily blended with the Paschen-$\rm{\gamma}$ and
C I lines to be identifiable.

A comparison of the cachito velocity, when interpreted as Si II, and the Fe II
velocity is plotted in Figure 14. We find excellent agreement between the two
velocities strengthening our conclusion that this feature is due to Si II
rather than high velocity hydrogen.

Figure 14: A comparison of the velocity of the cachito feature if it is Si II
($\rm{\lambda 6355}$) (black triangles) and the velocity of the metals in the
ejecta as characterized by the Fe II $\rm{\lambda 5169}$ lines (pink circles).
The velocities are the same indicating that this line is likely due to Si II
rather than high velocity hydrogen in the circumstellar medium.

The H $\rm{\alpha}$ emission profiles in the nebular spectra of ASASSN-15oz
are asymmetric (see Figure 15). In addition to emission at the rest wavelength
of H $\rm{\alpha}$, there is evidence of both blueshifted and redshifted
emission at 2200 km$\rm{s^{-1}}$ and 800 km$\rm{s^{-1}}$, respectively. This
feature is strongest in the first spectrum at day 228 (2016 Apr 11), becoming
weaker in the subsequent spectra. Asymmetric H $\rm{\alpha}$ emission has been
seen in a few IIP/IIL SNe (e.g. 1999em: Leonard et al. 2002; SN 2004dj: Chugai
et al. 2005; SN 2013ej: Utrobin & Chugai 2017). Leonard et al. (2002) explain
the shape of the H $\rm{\alpha}$ emission as either due to an asymmetric line-
emitting region or due to interaction between the ejecta and the CSM while
Chugai et al. (2005) and Utrobin & Chugai (2017) attribute this feature to a
bipolar $\rm{{}^{56}Ni}$ ejecta.

Figure 15: The asymmetric H $\rm{\alpha}$ emission in the nebular spectra of
ASASSN-15oz. The spectrum shows H $\rm{\alpha}$ emission at rest (solid gray
line), blueshifted by 2200 km $\rm{s^{-1}}$ and redshifted by 800 km
$\rm{s^{-1}}$ (dashed gray lines). Although the asymmetry is visible in all of
the spectra, it is most prominent in the earliest nebular spectrum on day 228
(2016 Apr 11) and could be a sign of interaction with CSM.

###  6.3 Radio

The origin of radio emission from SNe can be explained by the interaction of
the SN ejecta with the CSM. This interaction leads to a shockwave traveling
via the CSM that in turn accelerates electrons and enhances magnetic fields
and thus synchrotron emission ensues (Chevalier, 1982, 1998; Weiler et al.,
2002; Chevalier & Fransson, 2006). It has been shown over the last decade that
in most observed SNe, the optically thick part of the observed radio spectrum
can be modeled by a synchrotron self absorption (SSA; e.g. Soderberg et al.
2012; Horesh et al. 2013). We use the Chevalier (1998) SSA formalism to model
the radio measurements of ASASSN-15oz. In such a model the radio emission
peaks at a frequency below which the emission is strongly absorbed by SSA.
Above that frequency the emission becomes optically thin. As the shockwave
progress outwards to lower CSM densities, the SSA optical depth drops and the
peak of the radio emission moves to lower and lower frequencies. We use
Equation 1 from Chevalier (1998) to model the radio dataset (see Figure 16).
We use the standard equipartition assumption and set the microphysical
parameters, $\rm{\epsilon_{e}=\epsilon_{B}=0.1}$ (Chevalier & Fransson, 2006).
We also set the electron power-law distribution to $\rm{p=3}$, leaving the CSM
density and shockwave velocity as free parameters. We find the shockwave
velocity (assuming constant velocity) is $\rm{\sim 1.4\times 10^{4}}$ km
$\rm{s^{-1}}$ (which is the typical average value observed in SNe; e.g.
Chevalier & Fransson 2006). Assuming CSM surrounding the SN is created by a
stellar wind from the progenitor prior to explosion, we model the density
profile of the CSM as:

$\rho(r)=\frac{\dot{M}}{4\pi r^{2}v_{wind}}=\frac{K}{r^{2}}$ (3)

where $\rm{\dot{M}}$ is the wind mass-loss rate and $\rm{v_{wind}}$ is the
wind velocity. We find $\rm{K=4.51\times 10^{11}}$ g $\rm{cm^{-1}}$.

With a wind velocity between 10-100 km $\rm{s^{-1}}$, this implies a mass-loss
rate between $\rm{\dot{M}}$ $\rm{\approx 0.9-9\times 10^{-7}}$
M⊙$\rm{yr^{-1}}$.

Figure 16: The best fit SSA model (dotted lines) plotted over the observations
(symbols) at 22 GHz (black), 15 GHz (pink), and 4.8 GHz (indigo). These data
are best fit by a shock velocity of $\rm{\sim 1.4\times 10^{4}}$ km
$\rm{s^{-1}}$ and a density of $\rm{K=4.51\times 10^{11}}$ g $\rm{cm^{-1}}$

###  6.4 Light Curve Modeling with SNEC

We fit the light curve of ASASSN-15oz using the Supernova Explosion Code
(SNEC) (Morozova et al., 2015). SNEC is an open source Lagrangian 1D radiation
hydrodynamic code that employs flux-limited radiation diffusion and assumes
local thermodynamic equilibrium (LTE). SNEC uses a Paczyński equation of state
(Paczynski, 1983) solving for the ionization fractions using the Saha
equations in the non-degenerate approximation (Zaghloul et al., 2000).
Opacities are drawn from OPAL type II opacity tables (Iglesias & Rogers, 1996)
at high temperatures ($\rm{T=10^{4.5}-10^{8.7}}$ K) and tables of Ferguson et
al. (2005) at low temperatures ($\rm{T=10^{2.7}-10^{4.5}}$ K) supplemented
with an opacity floor of 0.01 cm2g-1. The mixing of isotopes due to the
explosion is modeled by boxcar smoothing the progenitor composition profile
with a kernel of 0.4 M⊙.

In the past, hydrodynamic SN models that used progenitors evolved with stellar
evolution codes (referred to as evolutionary progenitors for the remainder of
this paper; e.g. Kepler code, Weaver et al. 1978; Woosley & Heger 2007, 2015;
Sukhbold & Woosley 2014; Sukhbold et al. 2016; MESA, Paxton et al. 2018)
failed to reproduce the shape of the early observed light curves of Type II
SNe. Some authors have solved this issue by using hydrostatic progenitor
models with extended envelopes and density profiles that are set to match
observations rather than calculated by stellar evolution codes (referred to as
non-evolutionary progenitors for the remainder of this paper; e.g. Utrobin &
Chugai 2008, Utrobin & Chugai 2017). However, adding a dense CSM ejected just
before the explosion to progenitors evolved with stellar evolution codes has
resolved the discrepancy between the early observed light curves and the light
curves modeled using physically derived progenitors (e.g. Morozova et al.
2018; Gezari et al. 2015; Förster et al. 2018). The presence of this CSM is
corroborated by other observations of IIP/IIL SNe (e.g. Yaron et al. 2017;
Bullivant et al. 2018).

For the progenitor models, we use a set of non-rotating solar metallicity RSGs
parameterized by the $\rm{M_{ZAMS}}$, evolved with the Kepler code and
described in Sukhbold et al. (2016). Note that the stellar evolution
calculations made with Kepler take into account the regular steady winds
observed in RSGs, in the way prescribed by Nieuwenhuijzen & de Jager (1990)
and Wellstein & Langer (1999). For this reason, the final pre-SN masses of the
models may be up to several $\rm{M_{\odot}}$ smaller than their initial
$\rm{M_{ZAMS}}$. Both the $\rm{M_{ZAMS}}$s and pre-SN masses of the models
used are given in Table 2. However, for the typical mass loss rates of
$\rm{\lesssim 10^{-6}}$ M⊙$\rm{yr^{-1}}$, the density in these regular steady
winds is too low to have any noticeable effect on the post-explosion optical
light curves, and these winds are not included in the progenitor profiles.

We explore variations in CSM due to this enhanced mass-loss by adding a
steady-state wind above the RSG models with the density profile expressed in
Equation 3. This approximation allows us to parameterize the CSM with two
parameters, K and $\rm{R_{ext}}$, the radial extent of the CSM. In reality,
the late enhanced mass loss does not have to be in the form of a steady wind
with constant $\rm{\dot{M}}$ and may instead represent one or several eruptive
outbursts. However, in our experience, the early optical light curve depends
weakly on the exact slope of the CSM profile, being more sensitive to the
total CSM mass.

Recently, Paxton et al. (2018) demonstrated that IIP/IIL SNe are well
characterized by the radiation diffusion approximation from shock breakout
through the fall from plateau. At the same time, they have shown that the
assumption of LTE is not well satisfied at the photosphere location of the
IIP/IIL SNe models. For this reason, SNEC color light curves generally rise
faster than the light curves obtained from more sophisticated multi-group
radiation-hydrodynamics codes, like STELLA (Paxton et al., 2018). However,
this fact makes the case for introducing a CSM surrounding the RSG before its
explosion even stronger. In this view, the total CSM mass derived in our
analysis may be considered as a lower limit, while the values obtained with
more advanced codes should be comparable, or in some cases larger (Moriya et
al., 2017; Paxton et al., 2018).

To account for the formation of the neutron star, we excise the inner
$\rm{1.4\,M_{\odot}}$ of the progenitor models prior to the explosion. After
that, we model the explosion of ASASSN-15oz using a thermal bomb with
explosion energy $\rm{E_{exp}}$ lasting for one second in the inner 0.02 M⊙.
When computing the amount of energy injected in the form of a thermal bomb,
SNEC automatically takes into account the total initial (negative, mostly
gravitational) energy of the models. Therefore, by the explosion energy
$\rm{E_{exp}}$ in our analysis we mean the total energy of the models after
explosion, which is conserved in the code to better than 1% accuracy and
eventually mainly transformed into the kinetic energy of the expanding
envelope.

SNEC does not model nuclear reaction networks, but rather takes as input a
mass of 56Ni. We use a grid of 56Ni masses informed by our analysis in Section
3. While SNEC allows for the 56Ni to be mixed out to different values of mass
coordinate, Morozova et al. (2017) find that the progenitor masses and
explosion energies derived from fitting the IIP/IIL SN light curves are not
very sensitive to the degree of 56Ni mixing. Therefore, in our study we choose
to keep this parameter fixed and mix the 56Ni up to 5.0 M⊙.

We use SNEC to find the best progenitor parameters varying the progenitor
mass, explosion energy, 56Ni mass, CSM density, and CSM extent. Table 2 gives
the set of parameters used, resulting in over 5000 model light curves.

Table 2: The grid of parameters used by SNEC. The values that best fit the data are in bold. The pre-explosion masses are taken from Table 2 of Sukhbold et al. (2016). Parameter | Values
---|---
Progenitor ZAMS Mass ($\rm{M_{ZAMS}}$; M⊙) | 11, 13, 14, 16, 17, 18, 21
Final Pre-SN Mass ( M⊙) | 10.688, 11.567, 12.079, 13.145, 14.301, 14.936, 16.119
Pre-SN Radius (100 R⊙) | 5.7, 7.0, 7.8, 8.9, 9.1, 9.7, 11.2
Explosion Energy ($\rm{E_{exp}}$; 1051 ergs) | 0.5, 0.8, 1.1, 1.4, 1.7, 2.0
CSM Density (K; 1017 g $\rm{cm^{-1}}$) | 0, 10, 20, 30, 35, 40, 50, 60
CSM Extent ($\rm{R_{ext}}$; 100 R⊙) | 0, 15, 18, 21, 24, 27, 30, 33
Ni Mass (M⊙) |  0.08, 0.09, 0.11

Morozova et al. (2018) find the best fit model with and without CSM as a two
step process. First they modeling the second half of the light curve,
characterized by the s2 slope without CSM to determine the best progenitor
mass and explosion energy. Then, fixing the explosion energy and progenitor
mass, they explore the CSM parameter space. This is computationally less
intensive than modeling the full parameter space and allows them to explore a
finer grid of parameters. Given the complexity of the parameter space, we
choose to model the best light curve with CSM, exploring the full set of
parameters simultaneously.

For ease of comparison with observations, SNEC uses the photospheric
temperature at each time step to compute a blackbody spectrum, which it
combines with different filter throughputs to output a light curve in Sloan
filters u, g, r, and i, Bessell filters U, B, V, R, and I and PanSTARRS filter
z. During the rise and plateau phase, a blackbody should be a good
approximation to the longer wavelengths. However, line blanketing may cause
the bluer filters to be a poor representation of the observed spectrum (Kasen
& Woosley, 2009; Dessart & Hillier, 2005). For this reason the best fit model
is determined using the g, r, and i filters. While we do have a V-band light
curve, the throughput heavily overlaps with the g\- and r-bands and its
inclusion would give more weight to these wavelengths without providing new
information. The best fit model is determined by interpolating the well
sampled model to the observed wavelengths and computing a chi-square
minimization across all three filters. Given the uncertainty in the explosion
time, we shift the model spectrum by $\rm{\pm}$4 days and treat this offset as
a free parameter (toffset). Parameters for the best fit models with and
without CSM are given in Table 3. These models are shown with (solid lines)
and without (dashed lines) CSM interaction in the top panel of Figure 17.
Integrating the CSM density over it radial extent, we find a total CSM mass of
1.5 M⊙.

Table 3: The best fit SNEC model parameters for the simultaneous fits to g,r, and i-bands and to the bolometric luminosity with and without CSM. | Fit to g, r, i-bands | Fit to Bolometric Luminosity
---|---|---
Parameter | With CSM | Without CSM | With CSM | Without CSM
Progenitor ZAMS Mass ($\rm{M_{ZAMS}}$; M⊙) | 17 | 18 | 18 | 17
Final Pre-SN Mass ( M⊙) | 14.301 | 14.936 | 14.936 | 14.301
Pre-SN Radius (100 R⊙) | 9.1 | 9.7 | 9.7 | 9.1
Explosion Energy ($\rm{E_{exp}}$; 1051 ergs) | 1.4 | 2.0 | 1.4 | 2.0
CSM Density (K; 1017 g $\rm{cm^{-1}}$) | 40 | 0 | 60 | 0
CSM Extent ($\rm{R_{ext}}$; 100 R⊙) | 1800 | 0 | 1500 | 0
Ni Mass (M⊙) | 0.083 | 0.083 | 0.083 | 0.083
Total CSM Mass (M⊙) | 1.5 | 0 | 1.4 | 0

As a sanity check, we derive the bolometric luminosity by fitting a black body
to the photometric observations and find the best fit SNEC model with and
without CSM. The parameters for these models are given in Table 3 and are
plotted over the bolometric luminosity in the bottom panel of Figure 17. The
explosion energy and ejecta masses are similar to those derived from a
simultaneous fit of the g, r, and i bands. Although the individual CSM
parameters vary, these are highly degenerate and the total CSM mass, which our
model is a better measure of, is very similar. The difference in the total CSM
mass from the two different methods can be seen as an indication of the
uncertainty in this value.

Figure 17: Top: the multi-band optical light curve of ASASSN-15oz. The
observations are shown as circles and the best fit SNEC model is shown with
and without CSM (solid and dashed lines, respectively). We require CSM to fit
the light curve at early times. The blue filters are affected by line
blanketing (U and V-band) and are thus not expected to be well approximated by
the blackbody approximation used to convert the bolometric luminosity produced
by the modeling into magnitudes in individual filters. For this reason we do
not use these in the fit and have placed hatched lines over them. Line
blanketing may also be responsible for the overestimate of the model flux in
the g-band. Bottom: the bolometric luminosity of ASASSN-15oz plotted with the
best model fit to the bolometric luminosity including CSM (solid line) and the
best model without CSM (dashed line).

Towards a consistent picture of the SN explosion, we also verified that the
parameters we found from the light curve modeling were in agreement with the
interpretation of other observables. The high velocity derived from the
spectroscopy is consistent with the large mass and explosion energy found by
the light curve fitting. It is also compatible with the incomplete gamma ray
trapping, which is expected in SNe with high core velocities (Jerkstrand et
al., 2011). Figure 18 shows a comparison of the photospheric velocity of the
best fit SNEC model to the observed photospheric velocity of ASASSN-15oz.
Following Faran et al. (2014), we define the photospheric velocity as that of
Fe II ($\rm{\lambda}$5169). We use cachtio feature, now identified as Si II
($\rm{\lambda}$ 6355), to trace the photosphere at early times. We note that
while the Fe II ($\rm{\lambda}$5169) is often used to characterize the
photospheric velocity, it is not well understood where in the atmosphere line
is being formed. Paxton et al. (2018) find that the Fe II ($\rm{\lambda}$5169)
line is originating above the photosphere at a specific value of Sobolev
optical depth ($\rm{\tau_{sob}}$) (they find $\rm{\tau_{sob}\sim 1}$ fits the
observations well). However, Hamuy (2001), in comparing the velocity derived
from the Fe II ($\rm{\lambda}$5169) line to the photospheric velocity in the
Type II SN atmospheric models of Eastman et al. (1996), find a systematic
offset for individual objects that is in different directions for different
objects, leading to no average offset. While there is fair agreement between
the our model and observe velocity at early times, the model velocity is
significantly lower than the observed velocity starting around day 25.
Following Paxton et al. (2018), we adjust our model velocity to the value at
$\rm{\tau_{sob}=1.0}$. Figure 18 shows that this adjustment brings the SNEC
velocities closer to the observed velocities. A more detailed discussion
concerning the modeling of the minimum of the FeII line is beyond the scope of
this paper.

Figure 18: The model photospheric velocity computed by SNEC (indigo line)
compared to the observed photospheric velocity (symbols) taken to be the
velocity of the Si II ($\rm{\lambda}$ 6355) feature at early times (black
squares) and Fe II ($\rm{\lambda}$5169) feature at later times (pink circles).
The model deviates from the observations around the time that the photosphere
begins to recede into the ejecta. The model velocity at $\rm{\tau_{sob}=1}$ is
plotted in yellow. While this provides a better agreement with the
observations, the model velocity evolution is still faster than the observed
evolution.

According to the classical model by Chevalier (1982), the interaction of the
SN ejecta with the surrounding wind leads to the formation of a cool dense
shell between the forward and the reverse propagating shock waves. This shell
is responsible for the X-ray emission from Type II SNe, and it may explain the
broad base of the narrow emission lines seen in some interacting SNe (Chugai,
2001). In our SNEC model, the CSM density is too high and the density drop
between the envelope and the CSM is not sufficient to form a significantly
overdense shell. This is supported by Figure 19, where the interface between
the RSG envelope and the CSM corresponds to the velocity coordinate of
$\approx 4900\,{\rm km}\,{\rm s}^{-1}$. In this model, the forward shock wave
sweeps the CSM almost entirely before the breakout.

On the other hand, if we were to consider a three component model, consisting
of a RSG, a dense CSM and a regular low density stellar wind, like in Morozova
et al. (2018), the thin cool dense shell would be formed at the interface
between the dense CSM and the wind. In this scenario, the forward shock wave
keeps propagating into the low density wind and generates radio emission.
Since SNEC is not yet capable of handling the low density stellar wind, we
have omitted it in our current simulations. However, we already know from the
radio observations (Section 6.3) that the mass-loss rate of this wind in
ASASSN-15oz is $\approx 0.9-9\times 10^{-7}\,M_{\odot}\,{\rm yr}^{-1}$. Such
low density wind would not have any significant influence on the optical light
curve of the SN, which justifies using a two component model (RSG and dense
CSM only) for the purpose of this section.

Figure 19: Density as a function of velocity in our best fit SNEC model at the
free expansion stage (numbers in the plot indicate the time since explosion in
days). The interface between the RSG envelope and the CSM corresponds to the
velocity value of $\approx 4900\,{\rm km}\,{\rm s}^{-1}$. Mild increase in
density at $\approx 3500\,{\rm km}\,{\rm s}^{-1}$ is a result of the reverse
shock wave reflected from this interface.

##  7 Discussion

It is useful to compare observations of ASASSN-15oz to those of other IIP/IIL
SNe to understand if it is a typical IIL-like SN or a unique event. Although
they are part of a continuous class, IIL-like SNe typically have higher
velocity ejecta, are brighter, and have steeper light curve slopes than IIP-
like SN. In Section 5.1, we showed that ASASSN-15oz has above average ejecta
velocity. Here we examine the light curve slope and absolute magnitude in the
context of a sample of IIP/IIL SNe.

Population studies of IIP/IIL SNe have found that SNe with steeper light
curves are brighter (Li et al., 2011; Anderson et al., 2014; Sanders et al.,
2015; Valenti et al., 2016), due to a larger explosion energy, a larger
progenitor, CSM interaction, or a combined effect. The left panel of Figure 20
shows the V-band light curve of ASASSN-15oz compared to nine other well
studied SNe with a range of $\rm{s_{50V}}$. ASASSN-15oz is brighter and more
steeply declining than the classic IIP-like SNe 1999em and 2014et. We also
note that the slope of the radioactive tail of ASASSN-15oz is steeper than
other objects plotted here.

We use our database of publicly available light curves, SNDAVIS, to compare
ASASSN-15oz to a larger sample of objects. In the right panel of Figure 20 we
plot the $\rm{s_{50V}}$ and the absolute V-band magnitude for 105 SNe. As a
bright IIL-like SN, ASASSN-15oz is consistent with the sample correlation.

Figure 20: Left: the absolute V-band magnitude of ASASSN-15oz and nine well
studied SNe that span the IIP/IIL like decline rates (1979C: Balinskaia et al.
(1980); de Vaucouleurs et al. (1981); Barbon et al. (1982), 1980K: Barbon et
al. (1982); Buta (1982); Tsvetkov (1983), 1998S: Fassia et al. (2000); Liu et
al. (2000); Li et al. (2011); Pozzo et al. (2004), 1999em: Elmhamdi et al.
(2003), 2004et: Maguire et al. (2010), 2005cs: Pastorello et al. (2009),
2013by: Valenti et al. (2015), 2013ej: Valenti et al. (2014), 2014G: Terreran
et al. (2016)). ASASSN-15oz is both bright and steeply declining when compared
with this sample of classic SNe. Also visible in this plot, is the steeper
decline of the radioactive decay tail of ASASSN-15oz when compared with other
objects. Right: the absolute V-band magnitude compared to $\rm{s_{50V}}$ for a
sample of public SNe from the SNDAVIS database (black circles). The trend seen
by Anderson et al. (2014) and Valenti et al. (2016) of brighter SNe to have
steeper slopes is apparent. ASASSN-15oz (pink square) follows this
relationship. The SNe whose light curves are shown on the left are also
plotted on the right for reference. The gray dashed line denotes the
separation between IIP-like and IIL-like SNe. This separation is poorly
defined in the literature. The value we adopt is similar to that used by Faran
et al. (2014) and Li et al. (2011).

We also compare $\rm{s_{50V}}$ and the radioactive decay slope for ASASSN-15oz
with all SNe in our database for which both slopes are measured in Figure 21.
The steepness of $\rm{s_{50V}}$ in a IIP/IIL SN light curve can be explained
by the hydrogen recombination front receding faster than the ejecta expansion,
implying a progenitor with a smaller hydrogen envelope. The slope of the
radioactive decay tail is determined by the fraction of gamma rays trapped and
reprocessed by this same hydrogen envelope. A lower fraction of trapped gamma
rays can result either from a smaller hydrogen envelope or by external mixing
of $\rm{{}^{56}Ni}$. Given the possible dependence of both $\rm{s_{50V}}$ and
the radioactive decay slope on the hydrogen envelope, the correlation between
them, seen in Figure 21, is unsurprising. ASASSN-15oz, having relatively steep
slopes at both phases, lies at the IIL end of this relationship. This
correlation was presented in Anderson et al. (2014) and is supported here with
a larger sample.

Figure 21: Center: the slope of the light curve at late times, during the
radioactive decay phase compared to $\rm{s_{50V}}$ for a sample of publicly
available SN light curves. ASASSN-15oz is plotted in pink. Top (Right): A
histogram of the radioactive decay (photometric) phase slopes. All measured
slopes from this phase are included in the hatched sample. The objects for
which there is both a decay phase slope and $\rm{s_{50V}}$ are plotted in
black. The comparison of the full slope sample (hatch) to the combined slope
sample (black) gives us confidence that points we plot in the center are
representative of the full sample. The dashed grey line marks the radioactive
tail slope corresponding to complete trapping with a slope of
$\rm{s_{tail}=0.5}$ mag (50 day)-1. The dotted grey line represents the
approximate separation between IIP and IIL SNe with a slope of
$\rm{s_{50V}=0.5}$ mag (50 day)-1.

Having analyzed individual observations for CSM interaction, we now turn to
finding a unified model. We derived mass-loss parameters from both the light
curve modeling with SNEC and the modeling of radio observations. From the
optical light curve we found a compact CSM with a density of $\rm{K=4.0\times
10^{18}}$ g $\rm{cm^{-1}}$ extending to a radius of R = 1800 R⊙. Assuming a
typical of RSG wind velocity between 10-100 km $\rm{s^{-1}}$, the CSM density
required to fit the optical component implies a mass-loss rate between 0.8-8.0
M⊙$\rm{yr^{-1}}$. This is significantly higher than radio mass-loss rate,
$\rm{10^{-6}-10^{-7}}$ M⊙$\rm{yr^{-1}}$. While the mass-loss rate from the
radio observations is consistent with typical RSG mass-loss rates of
$\rm{10^{-4}}$\- $\rm{10^{-6}}$ M⊙$\rm{yr^{-1}}$, the optical mass-loss rate
is consistent with some late stage mass-loss scenarios (e.g. Chugai 2001;
Quataert & Shiode 2012; Shiode & Quataert 2014; Fuller 2017; Yaron et al.
2017). These discrepant values can be reconciled by considering the origin of
the radiation we are observing. The excess luminosity seen in the optical
light curve is due to radiation diffusion from the shocked CSM that is behind
the forward shock and is only visible after shock breakout. The radio
luminosity on the other hand, is created by the acceleration of electrons due
to the shock’s interaction with the CSM. It is therefore representative of the
shock’s location at the time of the observation. Our observations can then be
explained by a long period of average mass-loss, followed by a short period of
extreme mass-loss.

With this model in mind, we examine our UV and X-ray observations, in the
context of other IIP/IIL SNe. Strongly interacting SNe remain UV bright
throughout their interaction, implying that perhaps a UV excess would be
observed for IIP/IIL SNe during the first 30 days, when the CSM has the
greatest affect on the light curve. Furthermore, if in fact, the diversity in
light curve shapes is due to mass-loss and the interaction is strong enough to
produce a UV excess, then those SNe with a steeper slope, should have lost
more mass, and therefore should be more UV bright. We explore this by plotting
Swift UV color, UVW2 - V, over the first 30 days of evolution for eight SNe in
our database with Swift UVW2 and V-band photometry and $\rm{s_{50V}}$
measurements (Figure 22). The symbols for each SN are colored by the V-band
slope at 50 days. Although this sample is small, a wide range of slopes are
represented and no correlation between UV excess and slope is observed. We
interpret this as an indication that the CSM interaction is not strong enough
to produce UV excess. However, we cannot rule out the possibility that a
correlation is disrupted by uncertainties in the extinction correction.

Figure 22: A comparison of the UV color of ASASSN-15oz to that of all other
IIP/IIL SNe with Swift observations and $\rm{s_{50V}}$ measurements in
SNDAVIS. Each point is colored by the $\rm{s_{50V}}$ slope, with shallower
slopes in blue and steeper slopes in red. Based on this data, there is no
relationship between UV excess and slope.

Like the UV excess, X-rays from strongly interacting SNe are frequently
observed. To gain more insight into the X-ray upper limit, we compare the
upper limit of ASASSN-15oz with published X-ray observations of IIP/IIL SNe
found in the SNaX database (Ross & Dwarkadas, 2017) (see Figure 23). We find
our upper limit is above some other detections and upper limits of IIP/IIL SNe
in the database. This further suggests that deeper observations could be
needed to detect the X-ray flux of ASASSN-15oz.

Figure 23: The X-ray luminosity of all published SNe found in the SNaX
database. IIP-like SNe are shown with diamond symbols while IIL-like SNe are
shown with circular symbols and the upper limit of ASASSN-15oz is marked with
a square. Many detections from IIP/IIL SNe are below the upper limit derived
in this paper, implying that deeper observations would be required to detect
its X-ray flux.

##  8 Conclusions

We present extensive multi-wavelength observations of the type IIL-like SN,
ASASSN-15oz, including two epochs of radio observations, four epochs of NIR
spectroscopy, three epochs of NIR photometry, a 400 day dense multi-band
optical light curve, dense optical spectroscopy during the photospheric phase,
three epochs of nebular optical spectroscopy, one epoch of UV spectroscopy,
and 60 day UV and X-ray light curve.

Optical and NIR observations from the photospheric phase show ASASSN-15oz to
be a typical IIL-like SN, with a linearly declining light curve and above
average photospheric velocities. We find that in the explosion of ASASSN-15oz
0.08 -0.11 M⊙ of $\rm{{}^{56}Ni}$ was synthesized. However, since both the
light curve modeling and the nebular spectra point towards the lower limit, we
consider the lower limit most likely. Its bright peak magnitude is consistent
with the trend of increasing brightness with increasing photospheric slope.
The steep decline of the late time light curve is also congruous with the
direct correlation found between $\rm{s_{50V}}$ and radioactive tail slopes.
Optical light curve modeling using SNEC finds an 17 M⊙ progenitor, consistent
with the range of progenitors derived from modeling of the nebular spectra.

We search for signs of interaction at all wavelengths with the following
results:

  * •

We detect the ASASSN-15oz in radio observations, implying a mass-loss rate of
$\rm{0.9-9\times 10^{-7}}$ M⊙$\rm{yr^{-1}}$ for a wind velocity between 10-100
km $\rm{s^{-1}}$.

  * •

High velocity H and He are searched for in the NIR and optical spectra.
Although there is a feature bluewards of H $\rm{\alpha}$, no corresponding
feature is observed in any other hydrogen line nor in the He I ($\rm{\lambda
10830}$) line. For this reason we identify the cachito feature as Si II
($\rm{\lambda 6355}$) rather than high velocity hydrogen.

  * •

To accurately model the early optical light curve with an evolutionary
progenitor, we require CSM around the RSG progenitor with a density of
$\rm{4\times 10^{18}}$ g $\rm{cm^{-1}}$ and radial extent of 1800
$\rm{R_{\odot}}$, implying a short period of extreme mass-loss.

  * •

No narrow emission lines are detected in the optical and UV photospheric
spectroscopy, implying a mass-loss rate of $\rm{<10^{-4}}$ M⊙$\rm{yr^{-1}}$.

  * •

The H $\rm{\alpha}$ emission line profile in the nebular spectra is
asymmetric, which could be the result of interaction with CSM.

  * •

No UV excess is detected in the first 40 days of observations.

  * •

No X-ray signal is detected in the sum of all photospheric observations,
implying $\rm{L_{xray}<2.56\times 10^{39}}$ erg $\rm{s^{-1}}$.

Taking all of these results together, we find the progenitor of ASASSN-15oz
underwent standard RSG mass-loss for most of its evolution, with a short
period of extreme mass-loss just before explosion. It is possible that early
optical spectroscopy would have shown narrow emission lines as the shock
illuminated the edge of the dense shell of material. Additionally, an X-ray
detection may have been made with deeper observations. A dense shell of
material surrounded by a lower density CSM has now been used to model the
observations of a handful of objects (e.g. 1998S: Chugai 2001, 2013fs: Yaron
et al. 2017; 2013ej: Morozova & Stone 2018). As the best model to explain the
observations of a IIP-like SN (2013fs), a transition SN (2013ej), and now to a
IIL-like SN (ASASSN-15oz) it may be applicable to a broad range of
progenitors.

## Acknowlegments

This work was supported by the Swift Guest Observers Program through grant
NNX16AE90G. This work is based (in part) on observations collected at the
European Organization 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. Based on
observations collected at the European Organization for Astronomical Research
in the Southern Hemisphere under ESO program 095.A-0316(A). We acknowledge
support from EU/FP7-ERC grant [615929]. DAH, GH and CM are supported by NSF
grant AST-1313484. KML acknowledges funding from the European Research Council
under ERC Consolidator Grant agreement no 647208. A.L.P. acknowledges
financial support for this research from a Scialog award made by the Research
Corporation for Science Advancement. O.R. acknowledge support by projects
IC120009 “Millennium Institute of Astrophysics (MAS)” of the Iniciativa
Científica Milenio del Ministerio Economía, Fomento y Turismo de Chile and
CONICYT PAI/INDUSTRIA 79090016. DJS is a visiting astronomer at the Infrared
Telescope Facility, which is operated by the University of Hawaii under
contract NNH14CK55B with the National Aeronautics and Space Administration.
Research by DJS is supported by NSF grants AST-1821967, 1821987, 1813708 and
1813466. This research has made use of the NASA/IPAC Extragalactic Database
(NED), which is operated by the Jet Propulsion Laboratory, California
Institute of Technology, under contract with the National Aeronautics and
Space Administration. This work uses the National Radio Astronomy Observatory
which is a facility of the National Science Foundation operated under
cooperative agreement by Associated Universities, Inc. This work makes use of
observations from the LCO network. This work was partly supported by the UK
Space Agency. Based on observations obtained at the Gemini Observatory through
proposal GS-2016A-Q-75-25, acquired through the Gemini Observatory Archive,
and processed using the Gemini IRAF package, which is operated by the
Association of Universities for Research in Astronomy, Inc., under a
cooperative agreement with the NSF on behalf of the Gemini partnership: the
National Science Foundation (United States), the National Research Council
(Canada), CONICYT (Chile), Ministerio de Ciencia, Tecnología e Innovación
Productiva (Argentina), and Ministério da Ciência, Tecnologia e Inovação
(Brazil). This research made use of Astropy,999http://www.astropy.org a
community-developed core Python package for Astronomy (Astropy Collaboration
et al., 2013; Price-Whelan et al., 2018) Support for G.P. is provided by the
Ministry of Economy, Development, and Tourism’s Millennium Science Initiative
through grant IC120009, awarded to The Millennium Institute of Astrophysics,
MAS.

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##  Appendix A Tables of Observations

Table 4: Spectroscopic Observations of ASASSN-15oz.

Date | JD | Phase (Day) | Observatory | Instrument
---|---|---|---|---
2015-09-04 | 2457270.0 | 8.0 | LCO | FLOYDS
2015-09-05 | 2457270.7 | 8.7 | Swift | UVOTA
2015-09-05 | 2457271.1 | 9.1 | Swift | UVOTA
2015-09-05 | 2457270.6 | 8.6 | NTT | SOFI
2015-09-06 | 2457272.0 | 10.0 | LCO | FLOYDS
2015-09-07 | 2457273.0 | 11.0 | LCO | FLOYDS
2015-09-10 | 2457275.7 | 13.7 | Swift | UVOTA
2015-09-10 | 2457275.7 | 13.7 | Swift | UVOTA
2015-09-11 | 2457277.0 | 15.0 | LCO | FLOYDS
2015-09-11 | 2457276.9 | 14.9 | Swift | UVOTA
2015-09-16 | 2457282.0 | 20.0 | LCO | FLOYDS
2015-09-20 | 2457286.1 | 24.1 | LCO | FLOYDS
2015-09-21 | 2457287.5 | 25.5 | VLT | X-SHOOTER
2015-09-24 | 2457290.0 | 28.0 | LCO | FLOYDS
2015-09-30 | 2457296.0 | 34.0 | LCO | FLOYDS
2015-10-04 | 2457299.5 | 37.5 | NTT | EFOSC
2015-10-05 | 2457300.5 | 38.5 | NTT | SOFI
2015-10-06 | 2457301.9 | 39.9 | LCO | FLOYDS
2015-10-10 | 2457305.7 | 43.7 | IRTF | SpeX
2015-10-14 | 2457310.0 | 48.0 | LCO | FLOYDS
2015-10-25 | 2457320.9 | 58.9 | LCO | FLOYDS
2015-11-07 | 2457333.9 | 71.9 | LCO | FLOYDS
2015-11-08 | 2457334.5 | 72.5 | NTT | EFOSC
2015-11-19 | 2457345.5 | 83.5 | NTT | EFOSC
2016-04-11 | 2457489.9 | 227.9 | NTT | EFOSC
2016-04-11 | 2457489.9 | 227.9 | NTT | EFOSC
2016-06-09 | 2457548.8 | 286.8 | Gemini-S | GMOS
2016-06-10 | 2457549.7 | 287.7 | Gemini-S | GMOS
2016-06-12 | 2457551.7 | 289.7 | Gemini-S | GMOS
2016-08-03 | 2457603.7 | 341.7 | NTT | EFOSC
2016-09-11∗ | 2457642.6 | 380.6 | NTT | EFOSC
2016-09-11∗ | 2457642.6 | 380.6 | NTT | EFOSC
2016-09-19 | 2457650.5 | 388.5 | NTT | EFOSC
2017-09-20+ | 2458016.9 | 754.9 | Swift | UVOTA

∗No signal in data due to cloud cover
+ Template observation

Table 5: Sample of Imaging Observations of ASASSN-15oz. Full table available on-line. Date-Obs | JD | Phase (Day) | Apparent Magnitude | Apparent Magnitude Error | Filter | Source
---|---|---|---|---|---|---
2015-09-04 | 2457269.68 | 7.68 | 14.64 | 0.02 | B | LSC 1m
2015-09-04 | 2457269.68 | 7.68 | 14.66 | 0.02 | B | LSC 1m
2015-09-04 | 2457269.69 | 7.69 | 14.37 | 0.02 | V | LSC 1m
2015-09-04 | 2457269.69 | 7.69 | 14.43 | 0.01 | g | LSC 1m
2015-09-04 | 2457269.69 | 7.69 | 14.42 | 0.01 | g | LSC 1m

##  Appendix B Line Identification and Fitting Details

Table 6: The best parameters for the SYN++ fit. The Sobolev opacity $\rm{\tau}$ is modeled with an exponential profile with e-folding length, aux, minimum velocity, $\rm{v_{min}}$, and maximum velocity, $\rm{v_{max}}$. The temperature column is the Boltzmann excitation temperature. SYN++ models pure reasonance scattering which is a poor approximation for the H $\rm{\alpha}$. For this reason we use the H $\rm{\beta}$ line to determine the hydrogen contribution to the fit. A separate fit is performed for H $\rm{\alpha}$ and is listed separately in the table. We emphasize that this fit is used for line identification and not to derive ejecta properties. Ion | log($\rm{\tau}$) |  $\rm{v_{min}(kkm}$ $\rm{s^{-1})}$ | $\rm{v_{max}(kkm/s)}$ | aux | Temperature (kK)
---|---|---|---|---|---
H I -$\rm{\alpha}$ | 0.06 | 0.1 | 40.0 | 2.0 | 10.0
H I - $\rm{\beta}$ | 1.1 | 0.1 | 40.0 | 2.0 | 10.0
Na II | -0.5 | 0.1 | 40.0 | 1.0 | 10.0
O I | -0.8 | 0.1 | 40.0 | 1.0 | 10.0
Ca II | 1.3 | 0.1 | 40.0 | 2.0 | 10.0
Sc II | 0.1 | 0.1 | 40.0 | 1.0 | 10.0
Ti II | 0.3 | 0.1 | 40.0 | 1.0 | 10.0
Fe I | 0.3 | 0.1 | 40.0 | 1.0 | 10.0
Fe II | 0.3 | 0.1 | 40.0 | 1.0 | 7.0
Ba I | 0.0 | 0.1 | 40.0 | 1.0 | 10.0

For lines that are blended (but still showed partially isolated absorption
troughs for some features) we define the continuum as a straight line from the
lowest side of the feature. After dividing by the continuum, we then
simultaneously fit multiple 1D Gaussian profiles to the continuum normalized
spectra. For blended features originating in the same part of the ejecta, we
require that each component has the same width. We define the velocity of a
feature as the minimum (mean) of the individual Gaussian profiles
corresponding to that ion and the error as the standard deviation of the fit,
representing the range of velocities present in the photosphere. An example of
the multiple components and the combined fit is shown in the left panel of
Figure 24. For the Ca II NIR triplet we use the additional constraint of a
fixed offset between the minima of the features.

For unblended lines, we perform the following steps to find the minimum and
the standard deviation (following Silverman et al. 2012). We eliminate large
noise spikes (e.g. from cosmic rays) using a Savitzky-Golay smoothing filter
(Savitzky & Golay, 1964) with a quadratic function over a binsize of five
pixels. For each line, a minimum and maximum wavelength is defined as well as
a slope to account for the shifting of the feature over time as the ejecta
slows. A slope is then fit to small bins over this wavelength range, starting
with a binsize of five pixels and increasing until exactly three changes in
slope are found (at either edge of the line and the line center). If the
binsize reaches 40 percent of the feature size, no further attempt at fitting
is made. The feature edges are confirmed by fitting a quadratic function to
the region centered on each edge identified by the slope change with a width
of 20 pixels for the FLOYDS spectra and five pixels for the EFOSC2 spectra.
The edge is considered successfully found if the quadratic fit is concave
down. The continuum is defined by a line though these end points. Finally, the
velocity is found by fitting a cubic spline to the continuum subtracted flux
between the edges using the flux errors as weights. An example of a feature
with the continuum and spline fit is shown in the right panel of Figure 24.

Figure 24: An example of a fit to the multi-component Ca II NIR triplet (left)
using multiple Gaussians and a single line fit (following Silverman et al.
2012, right) to the Sc II ($\rm{\lambda}$5662) line. The observed spectrum is
plotted in black. The continuum edges are marked with a cyan circle and the
cyan dashed line connecting these points is used as the continuum. The Ca II
NIR triplet fit is found by simultaneously fitting 3 Gaussians with the same
standard deviation and mean offsets corresponding to the expected wavelength
separation of the triplet. The individual Gaussians are plotting at as dotted
lines and the combined fit is plotted as a solid green line. The minima of the
individual Gaussians is used to find the velocity of each component. The Sc II
feature is fit with a cubic spliine. The minimum of the spline is used to find
the Sc II velocity.

We use both methods to fit Na I ($\rm{\lambda}$ 5898) and H $\rm{\beta}$
($\rm{\lambda}$4861). For Na I, we find the spline fit does not characterize
the minimum of the profile well and instead use the Gaussian minimum. The H
$\rm{\beta}$ profile is contaminated by another feature (possibly Ti II) at
later times, offsetting the minimum found using the spline fit. For this
reason we prefer the Gaussian fit for H $\rm{\beta}$ as well. A comparison of
the results of both fits for H $\rm{\beta}$ at early times finds them in good
agreement, giving us confidence in the consistency of the two methods.