# Progenitor, environment, and modelling of the interacting transient, AT
2016jbu (Gaia16cfr)

S. J. Brennan1, M. Fraser1, J. Johansson2, A. Pastorello3, R. Kotak4H. F.
Stevance5, T. -W. Chen6,7, J. J. Eldridge5, S. Bose8,9, P. J. Brown10, E.
Callis1,R. Cartier11, M. Dennefeld12, Subo Dong13, P. Duffy1, N. Elias-
Rosa14,15, G. Hosseinzadeh16,E. Hsiao17, H. Kuncarayakti18,19, A. Martin-
Carrillo1, B. Monard20, G. Pignata21,22,D. Sand23, B. J. Shappee24, S. J.
Smartt25, B. E. Tucker26,27,28, L. Wyrzykowski29,H. Abbot26, S. Benetti3, J.
Bento26, S. Blondin30,31, Ping Chen28, A. Delgado32,33, L. Galbany34,M.
Gromadzki29, C. P. Gutiérrez19,35, L. Hanlon1, D. L. Harrison32,36, D.
Hiramatsu37,38,54,55,S. T. Hodgkin32, T. W.-S. Holoien39, D. A. Howell37,38,
C. Inserra40, E. Kankare4,S. Kozłowski29, T. E. Müller-Bravo41,53, K.
Maguire42, C. McCully37,38, P. Meintjes43,N. Morrell44, M. Nicholl45,46, D.
O’Neill25, P. Pietrukowicz29, R. Poleski29, J. L. Prieto22,47,A. Rau7, D. E.
Reichart48, T. Schweyer6,7, M. Shahbandeh49, J. Skowron29, J. Sollerman6,I.
Soszyński29, M. D. Stritzinger50, M. Szymański29, L. Tartaglia3, A.
Udalski29,K. Ulaczyk29,51, D. R. Young52, M. van Leeuwen32, B. van Soelen43
The authors’ affiliations are shown in Appendix A Contact e-mail:
<EMAIL_ADDRESS>
###### Abstract

We present the bolometric lightcurve, identification and analysis of the
progenitor candidate, and preliminary modelling of AT 2016jbu (Gaia16cfr). We
find a progenitor consistent with a $\sim$ 22–25 M⊙ yellow hypergiant
surrounded by a dusty circumstellar shell, in agreement with what has been
previously reported. We see evidence for significant photometric variability
in the progenitor, as well as strong H$\alpha$ emission consistent with pre-
existing circumstellar material. The age of the environment as well as the
resolved stellar population surrounding AT 2016jbu, support a progenitor age
of $>$10 Myr, consistent with a progenitor mass of $\sim$22 M⊙. A joint
analysis of the velocity evolution of AT 2016jbu, and the photospheric radius
inferred from the bolometric lightcurve shows the transient is consistent with
two successive outbursts/explosions. The first outburst ejected material with
velocity $\sim$650 km s-1, while the second, more energetic event, ejected
material at $\sim$4500 km s-1. Whether the latter is the core-collapse of the
progenitor remains uncertain. We place a limit on the ejected 56Ni mass of
$<$0.016M⊙. Using the BPASS code, we explore a wide range of possible
progenitor systems, and find that the majority of these are in binaries, some
of which are undergoing mass transfer or common envelope evolution immediately
prior to explosion. Finally, we use the SNEC code to demonstrate that the low-
energy explosion within some of these binary systems, together with sufficient
CSM, can reproduce the overall morphology of the lightcurve of AT 2016jbu.

###### keywords:

circumstellar matter – stars: massive – supernovae: general – supernovae:
individual: AT 2016jbu

††pubyear: 2021††pagerange: Progenitor, environment, and modelling of the
interacting transient, AT 2016jbu (Gaia16cfr)–References

##  1 Introduction

This is the second of two papers on the interacting transient AT 2016jbu
(Gaia16cfr). We report photometric and spectroscopic observations in Brennan
et al. 2021 (hereafter Paper I) and present an in-depth comparison of AT
2016jbu and SN 2009ip-like transients which include SN 2009ip (Fraser et al.,
2013a; Graham et al., 2014), SN 2015bh (Elias-Rosa et al., 2016; Thöne et al.,
2017), LSQ13zm (Tartaglia et al., 2016a), SN 2013gc (Reguitti et al., 2019)
and SN 2016bdu Pastorello et al. (2018). The work presented here will focus on
the progenitor candidate, its environment as well as modelling and
interpretation of the spectral and photometric evolution.

AT 2016jbu shows a smooth evolution of the H$\alpha$ emission profile,
changing from a P Cygni profile, typically seen in Type II supernova (SN)
spectra which show strong, singular peaked, hydrogen emission lines (Kiewe et
al., 2012; Taddia et al., 2015), to a double-peaked emission profile which
persists until late times, indicating complex, H-rich, circumstellar material
(CSM). AT 2016jbu and SN 2009ip-like objects show strong similarities in late
time spectra with strong Ca ii, He i and H emission lines as well as a lack of
any emission from explosively nucleosynthesised material such as
$[\text{O\,{i}}]~{}\lambda\lambda~{}6300,6364$ or $\text{Mg\,{i}}]~{}\lambda
4571$. No clear nebular phase is seen even after $\sim$ 1.5 years after
explosion in AT 2016jbu, and on-going interaction with CSM at late times may
be hiding a nebular phase and/or inner material from the progenitor.

The nature of SN 2009ip-like transients is much more contentious. On one hand,
there is evidence that these are genuine core-collapse supernovae (CCSNe), the
progenitor was destroyed and the transient will fade after CSM interaction
finishes (Pastorello et al., 2013; Smith et al., 2014a; Pastorello et al.,
2019a; Graham et al., 2014; Smith & Mauerhan, 2012). On the other hand, some
suggest these may be non-terminal events (Fraser et al., 2013a; Margutti et
al., 2014; Fraser et al., 2015; Graham et al., 2017), and SN 2009ip-like
events are a result of either pulsational-pair instabilities (Woosley et al.,
2007; Marchant et al., 2019), binary interaction (Kashi et al., 2013;
Pastorello et al., 2019a), merging of massive stars (Soker & Kashi, 2013) or
instabilities associated with rapid rotation close to the $\Omega\Gamma$-limit
(Maeder & Meynet, 2000).

As a follow-up to Paper I we continue the discussion on AT 2016jbu, focusing
on the progenitor and its local environment, as well as examining the
controversial topic of the powering mechanism behind SN 2009ip-like events. We
note that some of these topics have been discussed before by Kilpatrick et
al., 2018 (hereafter referred to as K18), and we refer to this work
throughout. For consistency with Paper I, and to compare to previous work by
K18, we take the distance modulus for NGC 2442 to be $31.60\pm 0.06$ mag. This
corresponds to a distance of $20.9\pm 0.58$ Mpc and adopt a redshift z=0.00489
from the H I Parkes All Sky Survey (HIPASS) (Wong et al., 2006). The
foreground extinction towards NGC 2442 is taken to be $A_{V}=0.556$ mag
(Schlafly & Finkbeiner, 2011) via the NASA Extragalactic Database
(NED;111https://ned.ipac.caltech.edu/). We correct for foreground extinction
using $R_{V}=3.1$ and the extinction law given by Cardelli et al., 1989. We do
not correct for any host galaxy or circumstellar extinction, however note that
the blue colours seen in the spectra of AT 2016jbu do not point towards
significant reddening by additional dust (further discussed in Sect. 4.2 and
Sect. 2). We take the V-band maximum at Event B (as determined through a
polynomial fit) as our reference epoch (MJD $57784.4\pm 0.5$; 2017 Jan 30).
Significant lightcurve features will use the same naming convention as in
Paper I for specific points in the lightcurve; Rise, Decline, Plateau, Knee,
Ankle.

In Sect. 5 we investigate the CSM environment around AT 2016jbu and using
photometry presented in Paper I reconstruct the bolometric evolution of Event
A and Event B up until the seasonal gap ($+140$ days), which we discuss in
Sect. 5.1. The progenitor of AT 2016jbu is discussed in Sect. 2 using pre-
explosion as well as late time imaging from the Hubble Space Telescope (HST).
This presence of pre-existing dust is discussed in AT 2016jbu using SED
fitting as well as dusty modelling in Sect. 3. Using HST and Very Large
Telescope (VLT) + Multi Unit Spectroscopic Explorer (MUSE) observations, we
investigate the surrounding stellar population and environment in Sect. 4. The
powering mechanism behind AT 2016jbu is discussed in Sect. 6. In Sect. 7.1,
the most likely progenitor for AT 2016jbu is examined. AT 2016jbu and most SN
2009ip-like transients display a high degree of asymmetry, most likely due to
a complex CSM environment, and this is expanded upon in Sect. 7.2. Finally, we
will address the explosion scenario for AT 2016jbu and perhaps other SN
2009ip-like transients, focusing on a CCSN scenario in Sect. 7.4, and an
explosion in a binary system in Sect. 7.5.

##  2 The progenitor of AT 2016jbu

The progenitor of AT 2016jbu was discussed by K18, who suggest that it was
consistent with an F8 type star of $\sim$18 M⊙ from an optical SED fit,
although circumstellar extinction places this as a lower bound.

There is a wealth of pre-explosion images of NGC 2442 and in this section we
explore this data to identify and characterise the progenitor of AT 2016jbu.
Here we are specifically concerned with the quiescent (or apparently
quiescent) progenitor which can only be identified in deep, high resolution
data.

###  2.1 Hubble Space Telescope imaging of the progenitor

Table 1: Observational log for all HST images covering the site of AT 2016jbu. Measured photometry (in the Vega-mag system) for AT 2016jbu is also reported. Phase is in rest frame days relative to Event B maximum light (MJD 57784.4)  Date | Phase (d) | Instrument | Filter | Exposure (s) | Mag (err)
---|---|---|---|---|---
2006-10-20 | $-3736.0$ | ACS/WFC | F435W | 4$\times$395 | 24.999 (0.037)
- | - | - | F658N | 3$\times$450 | 21.207 (0.024)
- | - | - | F814W | 3$\times$400 | 23.447 (0.019)
2016-01-21 | $-373.1$ | WFC3/UVIS | F350LP | 1$\times$420 | 23.625 (0.017)
- | - | WFC3/IR | F160W | 2$\times$503 | 20.726 (0.003)
2016-01-31 | $-362.9$ | WFC3/UVIS | F350LP | 2$\times$420 | 22.215 (0.026)
- | - | - | F555W | 2$\times$488 | 22.645 (0.002)
2016-02-08 | $-354.4$ | WFC3/UVIS | F350LP | 3$\times$420 | 22.134 (0.001)
- | - | WFC3/IR | F160W | 2$\times$503 | 19.570 (0.005)
2016-02-17 | $-345.5$ | WFC3/UVIS | F350LP | 3$\times$420 | 23.108 (0.012)
- | - | - | F814W | 2$\times$488 | 22.287 (0.003)
2016-02-23 | $-339.9$ | WFC3/UVIS | F350LP | 3$\times$420 | 23.212 (0.022)
2016-02-28 | $-334.8$ | WFC3/UVIS | F350LP | 1$\times$420 | 23.985 (0.022)
- | - | - | F555W | 2$\times$488 | 24.399 (0.004)
2016-03-04 | $-330.4$ | WFC3/UVIS | F350LP | 3$\times$420 | 22.729 (0.022)
- | - | WFC3/IR | F160W | 2$\times$503 | 20.224 (0.011)
2016-03-10 | $-323.8$ | WFC3/UVIS | F350LP | 3$\times$420 | 22.690 (0.037)
- | - | - | F814W | 2$\times$488 | 21.967 (0.022)
2016-03-15 | $-318.8$ | WFC3/UVIS | F350LP | 3$\times$420 | 22.868 (0.016)
- | - | WFC3/IR | F160W | 2$\times$503 | 20.323 (0.014)
2016-03-21 | $-313.1$ | WFC3/UVIS | F350LP | 3$\times$420 | 23.400 (0.013)
- | - | - | F555W | 2$\times$488 | 23.962 (0.012)
2016-03-30 | $-304.2$ | WFC3/UVIS | F350LP | 3$\times$420 | 23.775 (0.006)
- | - | WFC3/IR | F160W | 2$\times$503 | 21.301 (0.020)
2016-04-09 | $-293.7$ | WFC3/UVIS | F350LP | 3$\times$420 | 23.767 (0.006)
- | - | - | F814W | 1$\times$488 | 23.079 (0.035)
2019-03-21 | +776.7 | WFC3/UVIS1 | F555W | 320,390 | 23.882 (0.025)
- | - | - | F814W | 2$\times$390 | 23.239 (0.032)
2019-03-31 | +787.0 | ACS/WFC | F814W | 4$\times$614 | 23.529 (0.014)

NGC 2442 was observed with the Hubble Space Telescope (HST) on a number of
occasions both prior to and after the discovery of AT 2016jbu using the
Advanced Camera for Surveys (ACS) and both the UV-Visible and IR channels of
the Wide Field Camera 3 (WFC3/UVIS and WFC3/IR). We retrieved all images where
the image footprint covered the site of AT 2016jbu from the Mikulski Archive
for Space Telescopes (MAST222mastweb.stsci.edu/), these data are listed in
Table. 1. In all cases, science-ready reduced images were downloaded. With the
exception of the late-time ACS images taken in 2019, all analysis was
performed on frames that have been already corrected for charge transfer
efficiency losses at the pixel level (i.e. drc/flc files). For the 2019 ACS
images corrections for charge transfer efficiency were applied to the measured
photometry.

Figure 1: 2″$\times$2″ cutouts of all HST images centered on the progenitor
candidate for AT 2016jbu. Columns are ordered in wavelength from left to
right.

In order to locate a progenitor candidate for AT 2016jbu, we aligned the
F814W-filter image taken in 2017, when the transient was bright, to the
ACS+F814W image from 2006, approximately ten years prior to discovery. Using
20 point sources common to both frames and within 20″ of AT 2016jbu, we derive
a transformation between the pixel coordinates with an root mean square (rms)
scatter of only 12 milliarcseconds (pixel scale $\sim$ 0.05 ″/pixel). A bright
source is clearly visible at the position, and we identify this as the
progenitor candidate. The progenitor candidate is shown in Fig. 1, and is the
same source as was identified by K18.

We performed Point Spread Photometry (PSF) fitting on all HST images using the
November 2019 release of the dolphot package (Dolphin, 2000), with the
instrument-specific ACS and WFC3 modules. In all cases, we performed
photometry following the instrument-specific recommendations of the dolphot
handbook333http://americano.dolphinsim.com/dolphot/dolphot.pdf regarding
choice of aperture size. The WFC3 images were taken at two distinct pointings,
and each set were analysed separately, otherwise each contiguous set of
imaging with a particular instrument were photometered together, using a
single deep drizzled image as a reference frame for source detection.
Examination of the residual images after fitting and subtracting a PSF to
sources in the field revealed no systematic residuals, indicating satisfactory
fits in all cases. We show the HST photometry for AT 2016jbu in Fig. 2.

Figure 2: Foreground extinction corrected HST lightcurves of AT 2016jbu and
its progenitor are shown in the top panel. We also include a DECam r-band
detection at $-352$ d as a red filled circles with error bars. Error bars for
HST measurements are smaller than the point sizes. The horizontal line is to
guide the eye in comparing the late time ($\sim+2$ year) and pre-outburst
($\sim-10$ year) F814W magnitudes. We also plot the F350LP and F160W
lightcurves with a line to help guide the eye. Colour curves, corrected for
foreground reddening, are shown in the bottom panel. Colours as offset for
legibility by the amounts stated in the legend.

We find that the photometry reported by K18 is fainter than what we measure,
with a difference of $\sim 0.5$ mag in F350LP. We compared our measured F350LP
magnitudes and those of K18 to the values reported in the Hubble Source
Catalog (HSC; Whitmore et al., 2016). As the magnitudes reported in the HSC
are in the AB mag system, we applied the conversion from AB to Vega mag before
comparing to our photometry. The HSC F350LP magnitudes are consistent with
what we report here, and we also see the same variability for the progenitor
candidate. The cause of the difference between our photometry and that of K18
hence remains unknown.

We note that the broad-band photometry from HST is more than likely affected
by the strong emission in H$\alpha$. In Fig. 3 we show the throughput of the
HST filters compared to a late-phase spectrum of AT 2016jbu. The long-pass
F350LP filter will contain flux from H$\alpha$. Fortuitously, H$\alpha$ falls
in the low-throughput red wing of the F555W filter, where it will have
negligible effect. To verify this, we used synphot (Lim, 2020) to perform
synthetic F555W-filter photometry on the $+271$ d spectrum of AT 2016jbu, and
on the same spectrum where H$\alpha$ has been excised. The latter returns a
magnitude that is only 0.05 mag fainter than the former, and so the F555W
filter is not significantly affected by line emission.

Figure 3: HST filters used for pre-explosion lightcurve, compared to the
$+271$ d spectrum. Only F350LP covers the strong H$\alpha$ emission seen at
this epoch.

The progenitor is relatively red, bright, and shows significant variability
over timescales of $\sim$weeks. Correcting for foreground extinction, in 2006
the progenitor candidate had an absolute magnitude in F814W = $-8.46\pm 0.06$
and an $F435W-F814W$ colour of 1.13$\pm$0.04 mag. This colour is consistent
with a yellow hyper-giant (YHG), and corresponds to a blackbody temperature of
6500 K (Drilling & Landolt, 2000).

However, the narrow-band $F658N$ magnitude, which covers H$\alpha$, is much
brighter than would be expected. This indicates that even ten years before the
eruption or explosion of AT 2016jbu its progenitor was characterised by strong
H$\alpha$ emission.

In early 2016, between seven and ten months prior to the start of Event A, NGC
2442 was observed repeatedly with WFC3 in F350LP, F555W, F814W and F160W. This
dataset gives us a unique insight into the variability of the quiescent
progenitor prior to explosion. We see that even in quiescence (arbitrarily
defined as when the progenitor is fainter than mag$\sim$-10), the progenitor
displays strong variability. In particular in the best-sampled F350LP
lightcurve, the progenitor varies in brightness by 1.9 mag in only 20 days. As
discussed by K18, such rapid variability is hard to explain (although there is
some similarity to the fast variability seen in the pre-explosion lightcurve
of SN 2009ip; Pastorello et al., 2013). While it is impossible to know if the
variability is periodic on the basis of the short time coverage available for
AT 2016jbu, if it is periodic then the apparent period is around 45 days
(found via a low order polynomial fit to the F350LP lightcurve).

The variability seen in F350LP in early 2016 is also seen in other bands,
which appear to track the same overall pattern of brightening and fading. Fig.
2 shows the colour evolution of F350LP-F555W, F350LP-F814W, F350LP-F160W. In
all cases (with the exception of the earliest F350LP-F160W colour, which is
likely due to a spurious F350LP magnitude) we see a relatively minor colour
change over three months. In fact, it is possible that the apparent small
shift towards bluer colours is simply due to H$\alpha$ growing stronger, which
would cause the F350LP magnitude to appear brighter, rather than any change in
the continuum temperature.

At late times the progenitor candidate for AT 2016jbu is still present. In
2019, over two years since the epoch of maximum light, a source is found at
approximately the same F814W magnitude as was seen in 2006. It is unlikely
that this source is a compact cluster, as the pre-explosion photometric
variability can only be explained if a single star is contributing most of the
flux. Moreover, we compared the 2006 F814W and 2019 F814W images, and find
that the position of the source is consistent to within 17 mas between the two
epochs. This implies that the same source is likely dominating the emission at
both epochs, and if there is an underlying cluster it must be much fainter
than the progenitor source.

###  2.2 Physical properties of the progenitor

In order to determine the luminosity and effective temperature of the
progenitor of AT 2016jbu, we consider the WFC3 photometry taken in early 2016.
As a first step, we normalize out the variability seen over this period so
that we can build an SED from photometry taken in different filters at
different epochs. To do this, we fit a linear function to the colour curves of
our HST observations. We disregard the first epoch for the F350LP$-$F160W
colour (which is significantly redder than the other epochs); this measurement
is unreliable as the progenitor was affected by bad pixels in two of the three
individual exposures. We then use the fitted functions to interpolate or
extrapolate the magnitude of AT 2016jbu in F555W, F814W or F160W as necessary.
Finally, we shift the SEDs up or down in magnitude so that they all have the
same F814W magnitude as the 2006 value. The resulting normalised progenitor
SEDs can be seen in Fig. 4.

Figure 4: HST SEDs for AT 2016jbu based on the early 2016 WFC3 imaging are
shown in black. All SEDs have been shifted so that their F814W magnitudes
match, as discussed in the text. The F350LP filter magnitudes have not been
included in the SED as they are strongly affected by H$\alpha$ emission. We
also plot a number of SEDs derived from MARCS models. In the lower panel we
show the 15 M⊙ MARCS models appropriate to cool red supergiants. As this model
grid does not extend above 4500 K, we also plot a set of 5 M⊙ models with
slightly higher log(g) in the upper panel. All models have been shifted so
that they match the F814W filter magnitude of the progenitor, and we can see
that while the cooler models can match the NIR part of the SED, hotter
temperatures are required to match the optical.

In order to determine a progenitor temperature from the observed SED, we
compare to MARCS stellar atmosphere models (Gustafsson et al., 2008). We used
the pysynphot package to perform synthetic photometry on the surface fluxes of
the models and hence calculate their magnitude in each of the F555W, F814W,
and F160W filters. We shifted each model so that it matches the 2006 MW
extinction corrected F814W absolute magnitude of the progenitor. In the lower
panel of Fig. 4 we compare to the spherically-symmetric MARCS models for 15M⊙
red supergiants (RSGs; log(g)=0) at solar metallicity. While we can see that
the models provide a reasonable agreement, it is clear that the warmest model
(at 4500 K) is still too red to match the F555-F814W colour of the progenitor,
implying that the progenitor is hotter than this. Conversely, the 4000 K model
provides a good match to the F814W-F160W colours of the progenitor. As the 15
M⊙ super-giant models cover a relatively small temperature range, we also
explored the 5 M⊙ spherically symmetric MARCS models at log(g)=1.0 which span
a broader range (upper panel in Fig. 4). We find that a 5000 K model can
reproduce the optical colours of the progenitor, while the NIR is better
matched with a cooler 4000 K model.

While AT 2016jbu does not appear to suffer from high levels of circumstellar
extinction around maximum light, we cannot exclude the possibility that the
progenitor colours are caused by close-in CSM dust that was subsequently
destroyed. To explore this possibility, we used the dusty (Ivezic & Elitzur,
1997) code to calculate observed SEDs for a grid of progenitor models allowing
for different levels of CSM dust. dusty solves for radiation transport within
a dusty medium.

Since a dust-enshrouded progenitor could be hotter than the range of
temperatures covered by the MARCS model grid, we used the PHOENIX
models444http://phoenix.astro.physik.uni-goettingen.de (Husser et al., 2013)
as our input spectra. The PHOENIX models cover the temperature range from
6000–12000 K in 200 K increments, and have log(g) between 1 and 2 dex. MARCS
models covering a temperature range from 2600–7000 K in 100 K increments and
log(g) between 1 and 2 dex were also tested as input to dusty. These models
were then processed by dusty, assuming spherically symmetric dust comprised of
50 per cent silicates and 50 per cent amorphous carbon. The dust density
followed a $r^{-2}$ distribution, with a radial extent varying between 1.5 and
20 times the inner radius of the dust shell. The dust mass is parameterized in
terms of the optical depth in $V$-band, $\tau_{V}$, which varied between 0 and
5. We expect the dust temperature to be relatively hot (Foley et al., 2011;
Smith et al., 2013). We vary the dust temperature at the inner dust boundary
between 1250-2250 K. For each temperature and dust combination, we calculated
synthetic $F555W-F814W$ and $F814W-F160W$ colours, and compared to the
foreground extinction corrected colours of the AT 2016jbu progenitor. In Fig.
5 we plot all models that have colours within 0.1 mag of the progenitor.

We find that we are able to match the progenitor colours with models with
temperatures of between $10^{3.7}$ and $10^{3.9}$ K, for a circumstellar dust
shell with optical depth $\tau_{V}$ between 0.7 and 1.5, and a dust
temperature between 1500-2000K, in agreement what was seen in the environment
of SN 2009ip (Smith et al., 2013), as well as the SN Imposter, UGC2773-OT
(Foley et al., 2011). Additionally we find little influence of the radial
extent of the dust on matching models.

We calculated a luminosity for each of these models by integrating over its
spectrum, and find that the progenitor had a luminosity log$(L)$ between 5.1
and 5.3 dex (depending on temperature and extinction). Comparing to the BPASS
single star evolutionary tracks at Solar metallicity in Fig. 5, we find that
these correspond to approximately the luminosity of a 22–25 M⊙star as it
crosses the HR diagram to become a RSG. We plot the DUSTY models that match
our progenitor measurements in Fig.5.

Figure 5:  Hertzsprung–Russell (HR) diagram showing single star evolutionary
tracks from BPASS (Eldridge et al., 2017; Stanway & Eldridge, 2018). We
include SN 2009ip at log$(L/L_{\odot})=5.9$ and log$(T_{eff})=3.92$ (Smith et
al., 2010; Foley et al., 2011), as well as SN 2015bh (Boian & Groh, 2018),
IRC+10420 (Klochkova et al., 2016), and UGC2773-OT (Smith et al., 2015).
$\eta$ Car is plotted (red triangles) at several phases given in parentheses
(Prieto et al., 2014). We include the progenitor estimates for AT 2016jbu from
K18, in both the “low” and “high” states as green stars . We highlight the
Yellow-Void between 7000 K and 10000 K (de Jager, 1998) and include the output
of our dusty modelling for AT 2016jbu using PHOENIX models (multi-coloured
triangles) and MARCS models (multi-coloured squares). The colour of each point
corresponds to its optical depth ($\tau_{\nu}$) which is provided on the
colour bar on the right. We include an inset of the region around the
progenitor in the top left of the plot.

##  3 Evidence for Dust

Figure 6: SED fit of AT 2016jbu at -11 days before Event B maximum. Extinction
corrected photometry are grouped to 1 day bins and weighted averaged. Flux
errors are given as standard deviation of bins. Horizontal error bars
represent approximate filter band-pass. Hot blackbody is given in blue, the
cooler blackbody in red, and the black line is compound model. We note a
similarity the SED for SN 2009ip presented in Margutti et al. (2014).

We present a SED model fitted to our $-11$ day dataset in Fig. 6. We fit at
this phase as it has the broadest wavelength coverage without the need for
interpolation. We fit two blackbody models to the photometric points; one
representing a hot photosphere, and the second fitted to the IR excess seen in
$H$, $K$, $W1$, and $W2$. A single blackbody does not fit observations seen at
$-11$ d before maximum. Allowing for a second cooler blackbody at a larger
radius gives a model that fits the data well. This additional blackbody is
consistent with warm dusty material at a distance of 170 AU and a temperature
of $T_{\rm BB}\sim$1700 K. This material provides an additional luminosity of
$2.7\text{\times}{10}^{7}\text{\,}\mathrm{L}_{\odot}$. The hot blackbody has a
radius of 36 AU, a temperature of $T_{\rm BB}\sim$ 12000 K and has an
integrated luminosity of $1.3\text{\times}{10}^{8}\text{\,}\mathrm{L}_{\odot}$
and represents $R_{\rm BB}$ at this time. We find a dust mass of
$M_{dust}\approx 2.27\times 10^{-6}$ M⊙ (Using eq.1 from Foley et al., 2011).
In comparison Smith et al. (2014b) finds a lower dust mass of $(3-6)\times
10^{-7}$ M⊙ for SN 2009ip. Additionally we note a similarity to the SED for SN
2009ip presented in Margutti et al. (2014). The IR excess may be caused by
thermal radiation of pre-existing dust in the CSM re-heated by an eruption at
the beginning of Event B, i.e. an IR echo. We can compute the radius within
which any dust will be evaporated/vaporised at the phase of our SED fitting.
The radius of this dust-free cavity is given by:

$R_{c}=\sqrt{\frac{L_{SN}}{16~{}\pi~{}\sigma~{}T_{evap}^{4}~{}\langle
Q\rangle}}$ (1)

where $R_{c}$ is the cavity radius, $L_{SN}$ is the luminosity of the
transient, taken to be $1.3\text{\times}{10}^{8}\text{\,}\mathrm{L}_{\odot}$,
$\sigma$ is the Stefan-Boltzmann constant and $\langle~{}Q~{}\rangle$ is the
averaged value of the dust emissivity. Assuming radiation is absorbed with
efficiency $\sim$unity by the dust, we find a cavity radius of $\sim$ 245 AU
for graphite grains ($T_{evap}=1900~{}K$) and $\sim$ 400 AU for silicate
grains ($T_{evap}=1500~{}K$). Both values are significantly larger than what
we find from our warm blackbody radius ($\sim$ 170 AU). A dust destruction
radius larger than the blackbody radius of our putative warm dust component
appears at first glance to be inconsistent. To ameliorate this we suggest that
the dust may not be homogeneously distributed, and could be in either
optically clumps or an aspherical region that provides some shielding from
evaporation. Over time, we expect that the dust is further heated and
destroyed during the rise to Event B maximum. We find that by maximum
brightness that this additional blackbody component is no longer needed,
suggesting that the dust causing this NIR excess has been destroyed. As
discussed in Paper I, as well as K18, there are Spitzer+IRAC observations of
the progenitor site of AT 2016jbu, which show tentative detections in 2003 and
2018. Using 2003 Spitzer/IRAC and 2016 HST/F160W observations, K18 find fits
consistent with a compact dusty CSM component with mass
$M_{dust}\approx$7.7\text{\times}{10}^{-7}\text{\,}\mathrm{M}_{\odot}$$ at 72
AU. This may represent a dusty shell that is later seen as our 170 AU warm
blackbody. However, due to the time frame between Spitzer/HST observations
there are large uncertainties on dust parameters from K18. Fitting Spitzer
data only gives a slightly higher $M_{dust}$ value of $\sim$ $10^{-6}$ M⊙ at
120 AU. Due to the erratic variability seen in AT 2016jbu it is uncertain as
to whether these dust shells are the same, as AT 2016jbu may have a stratified
CSM environment resulting from successive outbursts. Although there is strong
evidence for pre-existing dust, we do not see any signature for newly formed
dust in the environment around AT 2016jbu (Meikle et al., 2007; Smith et al.,
2008; Smith, 2011). We see no NIR excess in late time $J$ and $K$ bands in
late time photometry nor an IR excess evident in spectra. Furthermore, there
is no blue-shift in the core emission component in H$\alpha$ (Paper I), which
is another indicator of newly formed dust.

##  4 The environment of AT 2016jbu

Along with direct detections of progenitors, analysis of the resolved stellar
population in the vicinity of a SN has also been used to infer the progenitor
age and hence initial mass (Gogarten et al., 2009; Maund, 2017; Williams et
al., 2018). An advantage to this technique is that it will not be affected by
any peculiar evolutionary history or variability of the progenitor that may
cause it to appear less or more massive than it truly is. On the other hand,
using the environment around a SN is an indirect proxy for the progenitor age,
and is predicated on the assumption that the local stellar population is
coeval. This method is also complicated by possible contamination from other
stellar populations from multiple star formation episodes.

###  4.1 Hubble Space Telescope imaging of the environment

In order to study the population in the vicinity of AT 2016jbu we require
sources to be matched between different filter images. While this is
straightforward for bright sources such as the progenitor of AT 2016jbu it is
more challenging for fainter or blended sources, especially when images have
different pixel scales or orientations. We hence re-ran the photometry on a
subset of the HST images (F435W, F658N and F814W from 2006 Oct 20; F350LP and
F555W from 2016 Jan 31), using a single drizzled ACS F814W image as the
reference image for all filters.

We chose a projected radius of 150 pc (1.48″) around AT 2016jbu as a
compromise between identifying sufficient stars to be able to constrain the
population age and ensuring we are still sampling a local population that is
plausibly coeval with the progenitor. We also create a less restrictive
catalog of sources within a projected distance of 300 pc from AT 2016jbu, as
well as a more limited catalog of sources within 50 pc. After applying cuts to
select only sources with a point source PSF, dolphot detects 84 sources at
S/N$>3$ within 150 pc of AT 2016jbu, and 255 sources within 300 pc.

Figure 7: Color magnitude diagram (CMD) of the stellar population around the
site of AT 2016jbu. We show three different colour combinations each with
PARSEC isochrones with population ages (solid colored lines) given in the
upper legend. Yellow squares are point sources within 150 pc and green
diamonds are within 300 pc, while sources within 50 pc of the progenitor are
plotted in red with error bars. The progenitor of AT 2016jbu from the early
2016 HST observations is given as a gold star in each panel. Figure 8:
Probability distribution of the age of the 150 pc stellar neighbourhood of AT
2016jbu found using AgeWizard. The 90 percent confidence interval is
highlighted in grey.

In Fig. 7 we compare our 50 pc, 150 pc and 300 pc populations to a set of
PARSEC isochrones555http://stev.oapd.inaf.it/cgi-bin/cmd (Marigo et al., 2017;
Bressan et al., 2012) in three different filter/colour combinations. We use
the most recent version of the PARSEC models (version 1.2S; Chen et al.,
2015), and for the purposes of the comparison we have applied our foreground
reddening and distance modulus to the PARSEC models.

The progenitor of AT 2016jbu clearly stands out from the local population,
both in terms of its bright apparent magnitude and unusual colours. The colour
of AT 2016jbu should not be compared to these isochrones; not only will the
$F350LP$ filter be strongly affected by H$\alpha$ emission, but as the various
filter combinations plotted do not come from contemporaneous data, the
variability seen in the progenitor will significantly affect the apparent
colour.

Turning to the 150 pc population, it is clear that no source is found to be
brighter than the 10 Myrs isochrone, constraining the population to be older
than this age. We find a similar result looking at the wider environment
within 300 pc of AT 2016jbu as well as the closer in population within 50 pc.

Using AgeWizard and BPASS models (Eldridge et al., 2017; Stevance et al.,
2020a, b), we obtain a probability distribution for the age of the resolved
stellar population within 150 pc around AT 2016jbu (see Fig. 8). The 90
percent confidence interval is found to be $15$–$200$ Myrs. Additionally, we
can ascertain that the neighbouring population of AT 2016jbu is older than 10
Myrs (5 Myrs) with over 95 (99.8) percent confidence.

Therefore, there is no evidence for a very young environment which would be
expected for a 80 (or even 150) M⊙ progenitor as proposed for SN 2009ip and
$\eta$ Car (Smith et al., 2010; Foley et al., 2011).

###  4.2 MUSE-ing on the local environment

We further investigate the nature of AT 2016jbu by looking at its local
environment in Integral Field Unit (IFU) data. AT 2016jbu was observed on 2017
Dec 2 ($+303$ d) using the VLT equipped with the MUSE instrument in Wide Field
Mode. The date cube was obtained as part of a survey of SN late-time spectra
in conjunction with the AMUSING survey of SN environments (Galbany et al.,
2016; Kuncarayakti et al., 2020). We downloaded the pre-calibrated data cubes
from the ESO archive and present our data analysis for the environment around
AT 2016jbu in Fig. 9.

Figure 9: IFU analysis of the environment of AT 2016jbu. The spectral cube was
corrected for Milky Way extinction and redshift. Observations were taken on
2017 Dec 2 ($+303$ d). Data is orientated such that north is up and east is to
the left. Included in each panel is a horizontal scale bar showing 500 pc. We
include a white light image (5000–7000Å) (top left), an extinction map (top
middle) based on Domínguez et al. (2013), a velocity field plot from H$\alpha$
corrected for recessional velocity (top right), star formation rate based on
Kennicutt (1998) (bottom left) and a metallicity map (bottom right) based on
Dopita et al. (2016). The location of AT 2016jbu is marked with a red circle
of radius 3 ″. We also include the location of SN 1999ga as a square to the
south west of AT 2016jbu. Data is not shown where EW $<$ 1 Å or within 3 ″ of
AT 2016jbu.

We fit for spectral features at each spaxel using a Gaussian emission profile
with a linear pseudo continuum over a small wavelength range. For measuring
the ratio of H$\alpha$ and H$\beta$ for the extinction map we constrain the
ratio of the two emission lines such that $H{\alpha}/H{\beta}~{}\leq~{}2.85$
(Case B recombination). To exclude the effects of AT 2016jbu on the analysis,
we exclude any pixel within 3 ″ of AT 2016jbu. We do not account for any
stellar absorption effects and as such, values here are lower limits. For
completeness we include the extracted spectrum of AT 2016jbu in Fig. 10.

Figure 10: Extracted spectrum of AT 2016jbu from VLT+MUSE. Spectrum was
extracted using a 1″ aperture at the transient position and corrected for
redshift and Galactic extinction. We marks strong emission features in red and
several forbidden transition lines are marked in blue. The transient appears
relatively blue even at $\sim$ +10 months, a possible sign of ongoing
interaction.

We show the extinction map across the field of view (FOV) using the method in
Domínguez et al. (2013), measured using the Balmer decrement. A proxy for the
star formation rate (SFR) is measured using LHα (Kennicutt, 1998). LHα was
corrected for extinction using the Balmer decrement (Vale Asari et al., 2020).
We also plot a metallicity map using the metallicity indicators given by
Dopita et al. (2016).

Fig. 9 does not include the core of the host galaxy, nor the southern arm. AT
2016jbu is located north of the southern distorted spiral arm of NGC 2442 and
is still clearly present in NGC 2442 almost a year after maximum as seen in
the white light image constructed from the datacube. The FOV
($1^{\prime}\times 1^{\prime}$) does however include the location of SN 1999ga
(Pastorello et al., 2009) as well as a luminous region in the center frame.
This “Super-Bubble” has been noted by previous authors (Pancoast et al.,
2010), and is seen in the irregular kinematic pattern seen in the center of
the FOV. Placing an age on this region is difficult, but it is likely to have
formed within the last 150–250 Myr (Mihos & Bothun, 1997). This is a
spherical-looking area within the diffuse region to the south-west of the
nuclear region, with a diameter of $\sim$ 1.7 kpc.

This Super-Bubble region is in the vicinity of both AT 2016jbu and SN 1999ga
(Ryder et al., 2001; Pastorello et al., 2008; Pancoast et al., 2010). This
region shows a high SFR and is bright in B-band, both signs of massive star
formation. High SFR is linked with a high SN rate (Botticella et al., 2012)
and it is a fair assumption that the general location of this Super-Bubble is
likely to host CCSNe, as is obvious from SN 1999ga.

The top middle panel in Fig. 9 maps the extinction across the FOV using the
Balmer Decrement (Domínguez et al., 2013). We find a value for the local
extinction ($E_{B-V}<0.45$) within 500 pc of AT 2016jbu with a similar value
seen across the FOV. The top right panel in Fig. 9 gives the velocity
dispersion across the FOV. The location of AT 2016jbu lies in an area moving
at $\sim-100$ km s-1 (image corrected for red-shift: z = 0.00489). The bottom
left panel shows a pseudo SFR based on the extinction corrected H$\alpha$
emission (Kennicutt, 1998). The figure shows two bright regions of star
formation, which is clear from the white light image. AT 2016jbu is situated
on the outskirts of a moderate star-forming region, $\sim 10^{-6}$
$\mathrm{M}_{\odot}~{}\mathrm{yr}^{-1}$. SN 1999ga lies on the edge of the
brighter star forming region. We include a metallicity map (bottom right
panel) following the metallicity indicators from Dopita et al. (2016). The
full FOV yields an approximately solar environment, with the median
metallicity across the field as 8.66 dex ($Z\approx 0.015$).

##  5 Bolometric evolution of AT 2016jbu

The bolometric lightcurve for AT 2016jbu is computed using $ugiz$, $UBVR$,
$JHK$, Gaia G, $W1$ and $W2$ from WISE, as well as Swift+UVOT UVW2, UVM2,
UVW1, U, B, and V. All calculations were carried out using
Superbol666https://github.com/mnicholl/superbol (Version 1.7; Nicholl, 2018).
Effective wavelengths were taken from Fukugita et al. (1996) and zeropoint
flux energies were taken from Tonry et al. (2018), while Superbol was modified
to also handle our WISE data. Extinction values in each filter were computed
using the York Extinction Solver (McCall, 2004). All magnitudes were converted
to $F_{\lambda}$, and interpolated where necessary to account for epochs
without specific filter coverage, taking $r$-band as the reference filter.
Black body fitting is performed for photometric bands that are centered on
$\lambda~{}>~{}3000$ Å to avoid the effects of strong line-blanketing. We also
obtain a pseudo-bolometric lightcurve by directly integrating $F_{\lambda}$
using the trapezoidal rule between 0.2 and 4.5 $\upmu$m (UVW2 to W2). We
present the results of our blackbody fitting in Fig. 11.

Figure 11:  Blackbody luminosity (top panel), temperature (middle panel), and
radius (lower panel) of AT 2016jbu calculated using SuperBol (Nicholl, 2018).
In the second and third panel, we include $T_{\rm BB}$ and $R_{\rm BB}$ fits
to our optical spectra (Paper I). We include approximate epochs where specific
H$\alpha$ features emerge in the $R_{\rm BB}$ panel, as discussed in the text.
The green shaded region shows the linear distance travelled by the slower
moving material, Ejecta 1, causing the P Cygni absorption. The blue shaded
region is the same for the faster moving material, Ejecta 2. The lower and
upper bounds for each band are bulk and max velocities respectively.

AT 2016jbu is an interacting transient showing strong emission lines.
Interpreting the blackbody evolution of photometry alone may be misleading,
due to the uncertainty as to whether the photometry is continuum-dominated or
line-dominated. For completeness we investigate blackbody fits from our
optical spectra. A black body function was fitted to the optical spectra
presented in Paper I while excluding strong emission features and only fitting
for $\lambda~{}>~{}3000$ Å. We find excellent agreement with the blackbody
evolution from photometric and spectroscopic data until $\sim$+125 d. After
this time, our observations become strongly line-dominated and blackbody
fitting becomes unreliable.

###  5.1 Radius and Kinematics

We show the blackbody luminosity ($L_{BB}$), radius ($R_{\rm BB}$), and
temperature ($T_{\rm BB}$) fits from Superbol in Fig. 11. The H emission for
AT 2016jbu shows two distinct absorption components (see Paper I). The first
component is seen in a P Cygni profile that is present up until $\sim$0.5
years after Event B maximum. The second component is present for $\sim$1 month
with respect to its first emergence, and suggests some absorbing, High
Velocity (HV) material. A similar feature has been seen in other SN impostors
(Tartaglia et al., 2016b, e.g.) and is common in SN 2009ip-like transients.
The presence of two regions of material with different velocities lends
credence to the idea that the interaction with some material during, or prior
to, Event A is the main source of energy input for Event B (Fraser et al.,
2013a; Thöne et al., 2017; Elias-Rosa et al., 2016; Benetti et al., 2016).

To explore this scenario further, we assume some optically thick material
causing the P Cygni absorption was ejected at $\sim$-90 d (first detection of
Event A, see Paper I), although this ejection may have occurred earlier.

Fitting a P Cygni absorption profile gives a maximum velocity of $\sim-850$ km
s-1, with a bulk velocity of $\sim-600$ km s-1 for the slower absorption
feature seen in the Balmer lines. We refer to this material as Ejecta 1. The
higher velocity absorption (which we refer to as Ejecta 2) has a maximum
velocity from the blue edge of the line of $\sim-10000$ km s-1 with the bulk
of the material at $\sim-4500$ km s-1. Using these velocities, we attempt to
constrain ejection/collision times.

The ejection epoch for the material causing the second high velocity
absorption component is open to debate. There is no evidence for this
additional absorption in optical spectra at $-24$ d and it is only seen on
$-15$ d. Under the presumption that we do not see this shell of material (i.e
Ejecta 2) until it interacts with the pre-existing material or until it is no
longer occulted by an existing photosphere, we find that a shell moving at
$\sim$4500 km s-1 for $\sim$3 days can reach the distance of $R_{\rm
BB}\sim$$0.1\text{\times}{10}^{15}\text{\,}\mathrm{c}\mathrm{m}$. We include
the distance travelled by Ejecta 2 in Fig. 11 as a blue band. We can constrain
the ejection date of this HV material to $\sim$21 days before maximum light
with the collision date (when Ejecta 2 catches up to Ejecta 1) at $\sim$19
days before maximum light.

We draw attention to the blackbody evolution over the period $-19$ d to $-13$
d. During this timeframe we see an inflection between the decline of Event A
and the rise of Event B. Although we have low-cadence coverage during Event A,
the distance travelled by Ejecta 1 follows $R_{\rm BB}$ quite well during
Event A. $R_{\rm BB}$ then contracts slightly beginning around $-30$ d to a
minimum at $-19$ d. At $\sim-19$ d $R_{\rm BB}$ increases at a velocity
similar to the velocity profile of Ejecta 2. This implies that the blackbody
radius now follows this material, which is likely Ejecta 2 with additional
material swept-up from Ejecta 1 and some CSM material. While this is
undoubtedly a simplified picture, it appear that AT 2016jbu is potentially
consistent with two successive eruptions (either non-terminal or a CCSN) where
the collision of ejecta powers the luminosity of Event B.

We initially find $T_{\rm BB}$ at $\sim 5700$ K, which is roughly constant up
until $\sim-30$ d. $T_{\rm BB}$ evolves exponentially from 6000 K at $-30$ d
to 12000 K at $-12$ d. After Event B maximum (marked as Decline in Fig 11),
$T_{\rm BB}$ cools to $\sim 5100$ K at the Knee epoch and slightly increase to
$\sim 6000$ K at the beginning of the Ankle epoch.

It is important to note that we see both components in spectra during the
first month of Event B. Additionally, the FWHM and velocity offset does not
significantly evolve during the first few months (see Paper I).

This is likely due to Ejecta 1 or the CSM or both being highly asymmetric. We
are motivated by the spectral evolution of the H$\alpha$ profile, the
evolution of $R_{\rm BB}$ and the degenerate appearance of the H$\alpha$
emission lines in SN 2009ip-like objects, see Paper I for further discussion.
If Ejecta 2 is spherically symmetric (e.g. possibly a CCSNe), some material of
Ejecta 2 would not interact with Ejecta 1 and expand freely along the lower
density regions.

We include labels indicating when certain spectral components appear in the
H$\alpha$ in Fig. 11, bottom panel. We see that the HV blue absorption feature
coincides with the evolution of Ejecta 2; this absorption is clearly seen at
$-18$ d and is detected until $+5$ d with fitting-model dependent tentative
detections up to $+10$ d. This second absorption component appears during the
rise in $R_{\rm BB}$ during Event B and vanishes when $R_{\rm BB}$ reaches its
maximum at $\sim$+7 d. At $\sim$+9 d, $R_{\rm BB}$ remains at a constant value
and we see the emergence of a broad, red shoulder emission in H$\alpha$ at
$\sim~{}1400$ km s-1, FWHM $\sim~{}4000$ km s-1. This may follow material
expanding at $\sim 1400$ km s-1, a receding photosphere or both. Several days
later the blue emission feature appears in H$\alpha$ and remains until late
times. At $+18$d this blue emission is centered at $\sim~{}-2400$ km s-1 with
FWHM $\sim$3800 km s-1.

Photons from the interaction site between Ejecta 2 and Ejecta 1/CSM may be
diffusing outwards at this epoch. We see that the red shoulder emission only
appears after $R_{\rm BB}$ reaches its maximum values, shortly followed by the
blue shoulder emission a week later. This leads to our conclusion that Ejecta
1 is partially asymmetric and when Ejecta 2 collides with it, Ejecta 1 is
partially engulfed. The interaction between these two shells then becomes
apparent at $\sim$+7 d when the asymmetric emission features are clearly seen
in H$\alpha$.

The $R_{\rm BB}$ peaks at $1.2\times 10^{15}$ cm, at $\sim 1$ week after Event
B maximum and remains roughly constant until the Knee phase. Thereafter, there
is a drop of $\sim-5\times 10^{13}$ cm/day until the beginning of the Ankle
phase. $R_{\rm BB}$ remains roughly constant at $\sim 0.3\times 10^{15}$ cm up
until the seasonal gap begins at $+140$ d. This epoch coincides with a
narrowing of both red and blue emission features and an increase in Equivalent
Width (EW) of both components. This may represent a time when opacities drop
significantly and there is less photon scattering. Using this collision
scenario as the dominant energy input for this transient, we will explore the
necessary energy budget in Sect. 6. Using the evolution of $R_{\rm BB}$ we can
better understand the nature of the explosion of AT 2016jbu, and we will
further discuss this in Sect. 7.2.

##  6 Powering AT 2016jbu

The nature of the energy input of AT 2016jbu and SN 2009ip-like transients is
debated. If AT 2016jbu is indeed a CCSN then this energy comes from an
imploding iron core and the early lightcurve is powered by the fast moving SN
ejecta material. Ejecta interacting with a dense CSM can power the lightcurve
for many years (see Fraser, 2020, and references therein). If the transient is
a CCSN, after the ejecta expands and cools, the late time lightcurve is
powered by the radioactive decay of 56Ni. We discuss the possible presence of
56Ni in Sect. 6.1.

If AT 2016jbu is a CCSN, then it is spectroscopically classed as Type IIn,
meaning we see strong signs of interaction with a dense, slow-moving CSM. We
discuss the energy input from ejecta/CSM interaction in Sect. 6.2.

###  6.1 56Ni mass

A product of CCSNe is explosively synthesised radioactive 56Ni, whose decay
can power the late time lightcurve of H-rich supernovae, after the hydrogen
ejecta have fully recombined and any additional interaction has stopped.
Anderson (2019) find that for H-rich, Type II SNe, the median value for the
amount of 56Ni synthesised is 0.032 M⊙. We show our attempt to fit for a
nickel decay tail in Fig. 12 (green dashed line). We find that the pseudo-
bolometric lightcurve shows a decay that is consistent with that of
radioactive nickel decay during the Ankle stage.

Determining an explosion epoch for AT 2016jbu is contentious. The transient is
clearly detected at -90 d in VLT+FORS2 imaging. We determine in Sect. 5.1 that
a second eruption (that may represent a genuine CCSN) occurred at $\sim$-21 d.
Using eq. 6 from Nadyozhin (2003), and taking the explosion epoch as
$\sim$-90d, we find a value of $M_{Ni}\leq 0.033$ M⊙ and taking $\sim-21$ d we
find a value of $M_{Ni}\leq 0.016$ M⊙. Following the arguments made in Sect.
5.1 we will take the latter explosion date as the more plausible motivated by
the apparent second eruption at $-21$ d ( this is the explosion epoch
typically assumed in the literature ), indicating a potential CCSN.

This limit on 56Ni is consistent with other SN 2009ip-like transients.
However, it is clear that during this time there is still on-going CSM-
interaction, as demonstrated from the multi-component H$\alpha$ profile in
Paper I, and as such, this value should be considered a conservative upper
limit, assuming any 56Ni is produced at all.

###  6.2 CSM-ejecta interaction

Figure 12: Pseudo-bolometric luminosity of AT 2016jbu using SuperBol (Nicholl,
2018). We include the luminosity shock function (Eq. 4; solid red line) and a
radioactive decay tail fit (green solid line). Both functions are extrapolated
until the end of observations (dashed lines). Both function are fitted to the
post-ankle stage and we include a zoom-in of this area in the top right. We
find a 56Ni mass of 0.016 M⊙ (assuming the SN explosion date as $-21$ d) and
$\dot{M}$ is 0.05 $\mathrm{M}_{\odot}~{}\mathrm{yr}^{-1}$for Eq. 4.

A previously explored scenario for the double-peaked lightcurve of SN 2009ip-
like objects is Event A represents a low energy eruption from the progenitor
star and Event B is powered by the interaction between the ejecta from this
eruption, and some pre-existing CSM that was ejected in the preceding years
(Mauerhan et al., 2013a; Fraser et al., 2013a; Thöne et al., 2017).

We measure the radiated energy released from Event A ($-90$ d to $-21$ d) as
$3.15\text{\times}{10}^{47}\text{\,}\mathrm{e}\mathrm{r}\mathrm{g}$ and the
energy from Event B ($-21$ d to $+450$ d) as $\sim$
$1.81\text{\times}{10}^{49}\text{\,}\mathrm{e}\mathrm{r}\mathrm{g}$. Fraser et
al. (2013a) find a similar value for SN 2009ip at $\sim$
$1.8\text{\times}{10}^{49}\text{\,}\mathrm{e}\mathrm{r}\mathrm{g}$.

If we assume that Event A is a symmetric explosion (similar to that proposed
in Mauerhan et al., 2014) we can approximate it using an Arnett Model (Arnett
& Chevalier, 1996). Taking the diffusion timescale for a photon to be
$t_{d}\approx\frac{L^{2}}{D}$, where D is a diffusion coefficient with
$D\approx\lambda c~{}=~{}\frac{c}{\rho~{}\kappa}$. Assuming that Event A
corresponds to the adiabatic expansion of a photosphere, and assuming
$L\approx R$, we can describe the diffusion timescale as:

$\tau_{d}=\left(\frac{3~{}\kappa~{}M_{ej}}{4~{}\pi~{}c~{}v_{ej}}\right)^{1/2},$
(2)

by substituting in $R\approx\tau_{d}\times v_{sh}$ and
$\rho\approx\frac{3~{}M_{ej}}{4~{}\pi~{}R^{3}}$, where $M_{ej}$ and $v_{ej}$
are the ejecta mass and velocity respectively, $\kappa$ is the opacity of the
ejecta, and $c$ is the speed of light. We take the rise time in r-band of
Event A to be similar to the diffusion time, and we get a value of $\sim$60
days. We assume the P Cygni minima follows this dense material ejected prior
to, or during, the beginning of Event A, as suggested by Thöne et al. (2017)
for SN 2015bh. Using Eq. 2 and taking $v_{ej}\approx 700$ km s-1 and assuming
a mean opacity of $\kappa$ = 0.34 $cm^{2}g^{-1}$ (assuming $e^{-}$ scattering
dominates in the H-rich ejecta) we find $M_{ej}$ for the Event A is $\sim$0.35
M⊙ giving a kinetic energy of $\sim$
$1.7\text{\times}{10}^{48}\text{\,}\mathrm{e}\mathrm{r}\mathrm{g}$.

This value is a factor 10 less than is required to power Event B. This is a
crude approximation as we invoke spherical symmetry. To fully investigate the
mass of Ejecta 1, detailed hydrodynamic simulations are needed (e.g. Vlasis et
al., 2016; Suzuki et al., 2019), which are beyond the work presented in this
paper.

Assuming free expansion, the constrained ejection times, and velocities for
our multiple shell models given in Sect. 5.1, the beginning of Event B
coincides with material from both shells being at the same location $R_{\rm
BB}\approx$ $0.7\text{\times}{10}^{15}\text{\,}\mathrm{c}\mathrm{m}$ (Fig.
11). This suggests that Event B is powered from the collision at $\sim-19$ d
of Ejecta 2 which interacts with the slower moving material ejected at the
beginning of Event A, (Ejecta 1).

It is difficult to measure the mass of Ejecta 2. If we assume that Event B is
powered solely by CSM interactions, we calculate that $M_{ej}~{}\sim~{}0.37$
M⊙ travelling at $\sim$ 5000 km s-1 can account for the energy seen, while
allowing for an extremely low porosity (or overlapping surface area between
the colliding material) of $10\%$. This value will change depending on the
opening angle of the interaction site, as explored in disk interaction models
(Vlasis et al., 2016; Suzuki et al., 2019; Kurfürst et al., 2020).

Even with this conservative estimate, our values of $M_{ej}$ are much lower
than those seen in CCSNe or $\eta$ Car. However, extremely low porosity (e.g.
1$\%$) would allow for a few M⊙ of ejected material if we assume no input to
the lightcurve from radioactive decay.

Although observed after peak luminosity, both SN 2013L and SN 2010jl showed a
plateau phase after maximum light (Ofek et al., 2014; Taddia et al., 2020).
This trend is discussed by Chevalier & Irwin (2011); SN ejecta interacting
with a dense mass loss region can form a plateau in luminosity lasting the
duration of the shock interaction, and ending when the entire interaction
material is shocked. As the photon mean free path increases with the geometric
expansion of the CSM, the innermost regions of the interaction are revealed.
This was suggested to explain the double-peaked spectral profiles of SN 2010ij
(Ofek et al., 2014), SN 2013L (Taddia et al., 2020), and iPTF14hls (Andrews &
Smith, 2018; Sollerman et al., 2019; Moriya et al., 2020) at late times. We
use the emergence of the blue emission feature and the decrease of the peak
velocity offset as a proxy for the shock front. We discuss the evolution of
this feature in Paper I. We fit a declining power law function to the peak
velocity of the blue emission from +20 to +120 days which is well fit by:

$v_{blue}(t)\approx(1375\pm 25)\times\left(\frac{t}{100d}\right)^{-0.40\pm
0.03}\mathrm{km~{}s^{-1}},$ (3)

Both red and blue emission components follow Eq. 3 well (the red component has
a different normalisation constant) up until the seasonal gap ($+140$ days).
After that both components maintain at a higher velocity and coast at
$\sim\pm$1300 km s-1 up until the end of our spectroscopic observations
($+575$ days), see Paper I. Under the assumption of steady-state mass loss,
the luminosity from CSM-shock interaction can be described by:

$L_{sh}=\epsilon~{}\frac{1}{2}~{}\frac{\dot{M}}{v_{wind}}~{}v_{ej}^{3},$ (4)

where $L_{sh}$ is the luminosity from CSM-ejecta interaction, $\epsilon$ is
the conversion efficiency from kinetic to thermal energy (taken to be 50%,
typical of Type IIn SNe (Smith, 2017)), $v_{ej}$ is the ejecta velocity, which
is set to Eq. 3, and $v_{wind}$ is the wind velocity. We fit Eq. 4 to our
bolometric lightcurve during the period from the Knee stage up until the
beginning of the seasonal gap. Fitting to this time-frame gives an upper limit
for $\dot{M}\approx 0.05$ $\mathrm{M}_{\odot}~{}\mathrm{yr}^{-1}$, if we
assume an LBV wind with $v_{wind}\approx 250$ km s-1 (we find a similar value
for $v_{wind}$ from our earliest H$\alpha$ profile). Setting $v_{wind}\approx
700$ km s-1, the value of the P Cygni minima, we obtain $\dot{M}\approx
0.14~{}$ $\mathrm{M}_{\odot}~{}\mathrm{yr}^{-1}$.

We base the above calculations on the assumption that the luminosity between
+70 d to +140 d is shock-CSM interaction dominated, with no other major
contributing energy source i.e. no major contribution from radioactive decay.
If AT 2016jbu is surrounded by a dense, disk-like CSM, the assumption that
this phase is interaction dominated is motivated by models (e.g. Fig. 11 from
Vlasis et al., 2016). These models show a similar lightcurve shape to AT
2016jbu, including a tail resembling radioactive 56Ni decay at $\sim$+80 days
past maximum brightness (these models assume no 56Ni). Symmetric ejecta and
disk interaction models show that the energy input at the Knee stage is
dominated by this ejecta-disk interaction. We will return to the possibility
of disk-like CSM in Sect.7.2.

After the seasonal gap ($+140$ days), the velocity of the red/blue emission
does not follow Eq. 3 and the bolometric luminosity does not follow Eq. 4. At
this point the lightcurve has increased in brightness, which is clearly seen
in Fig. 12. However by $\sim 400$ d, $L_{bol}$ fades below the extrapolated
value from Eq. 4.

After the seasonal gap both red and blue emission lines have similar FWHM,
$\sim$1500 km s-1, with the red emission having a slightly larger width, but
converging to the FWHM of the blue component at $\sim$400 d. If the red/blue
emission follows the shock interaction, this suggest an increased velocity of
the shock front. Conserving mass flux in the shock we have
$\rho_{1}~{}v_{1}=\rho_{2}~{}v_{2}$ where subscript 1,2 represent the post-
and pre- shock regions respectively. If the shock transverses to a lower
density CSM environment this can account for the increased velocity seen. This
might indicate that the shock has now reached a lower density environment,
perhaps created by the series of outbursts in the years prior. However, it is
not obvious how interaction with a less dense region of CSM would account for
the increased luminosity as well as the increased strength of He i emission
lines (also seen in SN 1996al; Benetti et al., 2016) at this time (Paper I).

##  7 Discussion

In the following section we will discuss the nature of AT 2016jbu. There is
much debate as to the nature SN 2009ip-like objects (Pastorello et al., 2008;
Smith & Mauerhan, 2012; Fraser et al., 2013a; Smith et al., 2014a; Margutti et
al., 2014; Graham et al., 2014; Pastorello et al., 2019a). Any scenario for AT
2016jbu or SN 2009ip-like transients needs to account for all of the following
points:

  1. 1:

Outbursts reaching an absolute magnitude of $M_{r}\sim-11\pm 2$ mag seen in
the historic lightcurve of the transient.

  2. 2:

A faint event, reaching an absolute magnitude of $M_{r}\sim-13\pm 2$ mag.

  3. 3:

An second event a few weeks later, reaching an absolute magnitude of
$M_{r}\sim-18.5\pm 0.5$ mag and ejecting material with velocities up to
$\sim~{}10,000$ km s-1.

  4. 4:

Ejected 56Ni mass of $\lesssim$ 0.02 M⊙.

  5. 5:

No directly observed synthesized material, either from explosive
nucleosynthesis or late-stage stellar evolution.

A possible addition to this list is double-peaked emission lines. This is seen
in the majority of SN 2009ip-like transients although, ironically, not SN
2009ip itself.

We address the probable progenitor in Sect. 7.1. Using our high cadence multi-
chromatic photometry presented in Paper I, and the bolometric evolution from
Sect. 5, we will present a likely explosion model and circumstellar (CS)
geometry for AT 2016jbu, that can be extrapolated to other SN 2009ip-like
transients in Sect. 7.2. We will discuss the validity of a CCSN scenario in
Sect. 7.3 and Sect. 7.4, and the possibility of the progenitor being in an
interacting binary system in Sect. 7.5.

###  7.1 The Progenitor of AT 2016jbu and SN 2009ip-like transients

The events of SN 2009ip-like transients may represent a critical step in the
late time evolution of massive stars. A dramatic increase in luminosity allows
for supper-Eddington winds and high mass-loss rates, however the mechanism
resulting in these outburst is unknown. Observations of shock features in the
Homunculus Nebula around $\eta$ Car may even point to explosive mass loss.
Furthermore, in the classical picture, LBVs should not be SN progenitors as
they have just transitioned to the He-core burning stage in their core

It is generally thought that SN 2009ip-like transients arise from very massive
stars (Foley et al., 2011; Pastorello et al., 2013; Fraser et al., 2013a;
Smith et al., 2014b; Fraser et al., 2015; Smith et al., 2016b; Elias-Rosa et
al., 2016; Pastorello et al., 2019a). The progenitor of SN 2009ip is thought
to be a 60–80 M⊙ LBV from pre-explosion images (Smith et al., 2010; Foley et
al., 2011). However, this was measured in a single band only, which may be
strongly affected by H$\alpha$ emission. As shown in Fig. 4, the bright
contribution of H$\alpha$ in F350LP will provide misleading SED fitting
results. While LBVs experience erratic mass loss as they undergo a short
transition from O-Type to the WR stars, AT 2016jbu appears to be too low mass
($\sim$22 M⊙) to be consistent with the SN 2009ip progenitor. We note that
this relatively low mass was found while taking into account the effect of
H$\alpha$ emission on the SED.

Our analysis on the progenitor mass for AT 2016jbu is the most secure for any
SN 2009ip-like transient in literature, as it is based on a broad optical to
NIR SED, as well as on the local neighbourhood. From our SED fitting to the
early 2016 HST data, we find the color of the progenitor is consistent with a
yellow hyper-giant. Using dusty modelling and matching the output spectra to
these colour values we find values for L and T which are consistent with a
mass of single star of 22–25 M⊙, consistent with the results from K18.
Moreover, the local environment, which can be be assumed to be composed of a
similar stellar population, demonstrates that we can effectively rule out a
very young population (expected for a 60–80 M⊙ star).

In order to explore the progenitor further, we turn to a grid of stellar
models created with the BPASS code. The BPASS stellar model library contains
the time varying properties of over 250,000 star systems for a grid of initial
parameters and a population containing a realistic fraction of binary and
single star systems (Eldridge et al., 2017; Stanway & Eldridge, 2018). Using
hoki777https://github.com/HeloiseS/hoki (Stevance et al., 2020a), we searched
for models matching the observed temperature and luminosity of the progenitor
of AT 2016jbu, considering both the possibility of a terminal core-collapse
supernova and of a non-terminal event.

For the CCSN (non-terminal explosion) scenario we find 12 (1668) matching
stellar models, and 0 (3) of these models correspond to single star systems.

The ZAMS and final mass distributions, as well as the evolutionary tracks for
both interpretations, are presented in Fig. 13. We can evaluate the mean and
standard deviation for the two scenarios: $\overline{{\rm M}}=12.3~{}$M⊙,
$\sigma=1.9$ M⊙ and $\overline{{\rm M}}=22$ M⊙, $\sigma=3.4$ M⊙ for CCSN and
non-terminal explosion cases, respectively.

Figure 13: Mass distributions (ZAMS and final) and evolutionary tracks of
BPASS models matching the L and T derived from dusty modelling. Upper (lower)
panel shows the CCSN (Non-terminal) scenario. Each evolutionary track is
plotted at a low transparency and therefore the lighter the tracks, the rarer
they are in our matches. We mark the search area in T and L from Sect. 2 in
each HR diagram.

We find a mean lifetime of $7.3^{+0.1}_{-.1}$ Myr and $7.0^{+0.1}_{-0.1}$ Myr
for the CCSN and non-terminal explosion scenarios, respectively. Very massive
stellar progenitors (e.g. classical LBVs with $>50$ M⊙) are confidently
excluded for AT 2016jbu.

There are numerous suggestions in the literature that LBVs can be the direct
progenitors of CCSNe (e.g. Trundle et al., 2008; Dwarkadas, 2011; Smith &
Tombleson, 2015; Humphreys et al., 2016; Ustamujic et al., 2021). It has been
suggested that the LBV phenomenon may occur in stars with initial masses as
low as $20-25$ M⊙, particularly when rotation is included in models (Groh et
al., 2013). Such LBVs may appear similar to F-type yellow super giants during
their eruptive stage (Kilpatrick et al., 2018; Humphreys et al., 2016). So, it
is possible that the progenitor of AT 2016jbu is a low-mass LBV. However, we
still require a high mass-loss of $\sim$0.05
$\mathrm{M}_{\odot}~{}\mathrm{yr}^{-1}$ to explain the lightcurve of AT
2016jbu. This is not dissimilar to the mass loss rate of $\eta$ Car during its
Great Eruption ($\sim 0.1$ $\mathrm{M}_{\odot}~{}\mathrm{yr}^{-1}$; Davidson &
Humphreys, 1997), but it remains unknown whether lower mass LBVs can sustain
such a high mass loss rate.

###  7.2 Geometry of AT 2016jbu and SN 2009ip-like transients

An interesting problem to solve with the CCSN scenario is that of the presence
and geometry of the CSM, as discussed in Sect. 5. The LBV-type winds invoked
in Sect. 6.2 do not apply to lower mass progenitors; indeed we find an average
mass-loss rate over the last 1 Myr of $\log(\dot{\rm M})=-5.4^{+0.2}_{-0.8}$
$\mathrm{M}_{\odot}~{}\mathrm{yr}^{-1}$ and $\log(\dot{\rm
M})=-4.9^{+0.2}_{-0.3}$ $\mathrm{M}_{\odot}~{}\mathrm{yr}^{-1}$ for the CCSN
and non-terminal scenario, respectively.

One can sustain a dense CSM even with a low mass loss rate provided the wind
velocity is sufficiently small. Using $\log(\dot{\rm M})$/$v_{wind}$ as a
proxy for wind density, we compare the average ratio found in our models to
that assumed in Sect 6.2. We find that for both sets of progenitor models
$\dot{\rm M}$/$v_{wind}\approx 10^{-6}$, compared to a value of $\approx
10^{-4}$ found for AT 2016jbu. Thus, we can confidently assert that steady
winds are not able to create the CSM observed in AT 2016jbu.

The alternative is episodic mass loss resulting from Roche lobe overflow
(RLOF) or common envelope evolution (CEE). We examined the CCSN progenitor
models found in BPASS and find that 3 models are in a CEE phase at the time of
CCSN explosion; furthermore we find another 2 undergoing mass transfer.
Similarly, for the non-terminal models we find that 937 models are in the CEE
phase and 501 are undergoing stable mass transfer at the point where they
match the observed L and T of the AT 2016jbu progenitor. Consequently, the
BPASS models reveal that the peculiar combination of properties and
environment of AT 2016jbu can be explained by binary interactions.

A radially confined, dense, disk-like CS environment has been suggested for SN
2009ip-like transients (Smith et al., 2014a; Margutti et al., 2014; Levesque
et al., 2014; Margutti et al., 2014; Fraser et al., 2015; Benetti et al.,
2016; Tartaglia et al., 2016a; Pastorello et al., 2018; Andrews & Smith, 2018)
as well as other Type IIn SNe (van Dyk et al., 1993; Benetti, 2000;
Stritzinger et al., 2012; Benetti et al., 2016; Andrews et al., 2017; Nyholm
et al., 2017) and super-luminous supernovae (SLSNe) (Metzger, 2010; Vlasis et
al., 2016).

Double-peaked line profiles are signs of an asymmetric environments such as a
disk, rings, or bipolar outflows cause by an asymmetric explosion. This is
similar to the presence of double-peaked H$\alpha$ (and other emission lines)
originating from an accretion disks in active galactic nuclei (e.g.
Shapovalova et al., 2004) as well as double-peaked emission from Be/shell
stars (e.g. Andrillat et al., 1986), although their formation and powering
mechanism are extremely different. We show in Paper I that AT 2016jbu and
other SN 2009ip-like objects show a degree of degeneracy in the appearance of
their H$\alpha$ profiles, which may be explained with a simple viewing angle
effect.

We suggest that AT 2016jbu has underwent a series of eruptions, such as has
been suggested for $\eta$ Car (see review by Smith, 2009) and SN 2009ip
(Mauerhan et al., 2014; Levesque et al., 2014; Margutti et al., 2014; Reilly
et al., 2017), and a significant portion, if not all, of the explosion energy
is a result of a ejecta-ejecta or ejecta-CSM interaction, which dominates
around a month after maximum light. It is uncertain whether any of these
eruptions emanate from core-collapse.

Recently, several groups have modelled the interaction of ejecta with
aspherical CSM (Vlasis et al., 2016; McDowell et al., 2018; Suzuki et al.,
2019; Kurfürst et al., 2020; Nagao et al., 2020). Vlasis et al. (2016) has
modeled the lightcurve evolution of a spherically symmetric eject colliding
with a disk-like CSM. We find that similarities between these models and AT
2016jbu. One important feature is after $\sim$ +80days these models seem to
follow a decay similar to that expected from 56Ni. The energy source at this
time is solely powered from CSM interaction and not from radioactive decay.
However, these models cannot explain the increased brightness in AT 2016jbu
after the seasonal gap, although this likely reflects a clumpy CSM and would
require fine-tuning of the CSM density profile.

Models by Kurfürst et al. (2020) have modeled ejecta interaction with
aspherical CSM for a range of viewing angles (Model A and Fig. 12 in Kurfürst
et al., 2020) demonstrating a clear viewing angle degeneracy, with looking
down through the plane of the CSM showing the greatest “double-peaked”-ness
and looking through the material showing the least (i.e singularly peaked
emission lines). This can naturally explain the variations in H$\alpha$
appearance found amongst SN 2009ip-like transients (see Paper I).

For SN 2009ip-like transients, there is some discrepancy as to the eruption
epoch of this asymmetric structure, with some authors suggesting this material
was ejected close to/during Event A (e.g. Margutti et al., 2014; Tartaglia et
al., 2016b; Thöne et al., 2017) whereas some authors speculative the disk has
been ejected much longer before (e.g. Mauerhan et al., 2013b, 2014). This is
difficult to understand without specific stellar evolutionary models.

As discussed in Sect. 5.1, we proposed a double eruption model, where an first
ejecta interacts with pre-existing CSM, followed by a second eruption some
months later. The collision of these two ejecta produce the spectral and
lightcurve evolution we present in Paper I and can be extrapolated to fit the
observables of several SN 2009ip-like transients. We provide an illustration
in Fig. 14 with a detailed outline of events in the provided caption.

Figure 14: Toy model depicting possible geometry and explosion scenario for AT
2016jbu. This diagram illustrates the discussion in Sect. 5.1 and Sect. 6.2.
Left panels show a simplified illustration of the CS environment around the
progenitor at specific epochs. Middle column shows H$\alpha$ profile at
corresponding epochs. The upper right panel shows the bolometric luminosity
and the lower right shows the blackbody radius. We include the distance
travelled by Ejecta 1 (green shaded region) and Ejecta 2 (blue shaded region).
Event A begins with an eruption of ejecta 1 which originates from the
progenitor system. The eruption and expansion of Ejecta 1 is seen in $L_{bol}$
and $R_{\rm BB}$, both peaking at $\sim-27$ d. A dense disk-like CSM funnels,
and is partially engulfed by, Ejecta 1. At $\sim-21$ d, Ejecta 2 is ejected
with a velocity of $\sim 5000$km s-1(this could be the SN explosion) and
almost immediately collides with Ejecta 1 wit some fast moving material
escapes along less dense polar regions. $L_{bol}$ and $R_{\rm BB}$ follow the
expansion of an opaque Ejecta 2 following the HV material seen in H$\alpha$.
Ejecta 2 becomes optically thin and the photopshere begins to move inwards in
velocity space. There is a linear decay in $R_{\rm BB}$ until $\sim+22$ d or
the beginning of the plateau stage.$R_{\rm BB}$ plateaus at $\sim+25$ d due to
effective internal heating from the site of interaction. Photons originating
from the interaction site between Ejecta 1, Ejecta 2 and the CSM begin to
diffuse outwards, as the material becomes partially transparent. This
coincides with the metamorphosis of the blue HV absorption to an emission
profile. At $\sim+45$ d, the Knee stage drops in $L_{bol}$ and $R_{\rm BB}$
with $R_{\rm BB}$ at a slightly higher value when compared to the beginning of
Event B. Both red and blue emission lines narrow at this stage which may
signify any intervening material is now completely optically thin and any
escaping photons undergo minimal scattering. The dominant source of energy is
now shock interaction due to ejecta-csm interaction. In the bottom left panel
we also include a line of sight dependency. We expect the transient to show
double-peaked emission lines when observed near the equator (e.g. AT 2016jbu,
SN 2015bh, SN 1996al) and more singularly peaked when observed towards the
polar regions (e.g. SN 2009ip). We note the similarities between this toy
model and those proposed for $\eta$ Car (e.g. Smith et al., 2018).

###  7.3 Modelling the lightcurve using SNEC

To further explore the plausibility of the progenitor matching BPASS models
from Sect. 7.1, we exploded a small subset of these with the SuperNova
Explosion Code (SNEC, Morozova et al., 2015). The full details of how BPASS
models are exported and exploded within SNEC can be found in Eldridge et al.
(2019). The key addition to using the progenitor model structure is to add on
a CSM component around the star. Here we use the values derived in Sect. 5.1
of a terminal wind velocity of 250 km s-1and a mass-loss rate of 0.05
$\mathrm{M}_{\odot}~{}\mathrm{yr}^{-1}$. For each of the input stellar models
we use an explosion energy of $5.6\times 10^{49}$ ergs, 0.016 M⊙ of 56Ni, and
an inner mass cut at 5 M⊙ with nickel mixing out to 0.6 M⊙. The resultant
simulated bolometric lightcurves are shown in Fig. 15 and model parameters are
given in Table 2.

Figure 15: Diagram showing the observed lightcurve and lightcurves simulated by SNEC from progenitors which match the pre-explosion constraints. All include the circumstellar medium as described earlier. We include our pseudo-bolometric for AT 2016jbu in black and the optical pseudo-bolometric lightcurve for SN 2009ip in blue(Fraser et al., 2013b; Pastorello et al., 2013). Table 2: The parameters of the BPASS models exploded with SNEC. $M_{1}$ | $M_{2}$ |  | $M_{\rm final}$ | $M_{\rm CO}$
---|---|---|---|---
$/M_{\odot}$ | $/M_{\odot}$ | $\log(P_{\rm i}/{\rm days})$ | $/M_{\odot}$ | $/M_{\odot}$
17 | 11.9 | 3 | 5.9 | 4.0
18 | 16.2 | 3 | 6.5 | 4.1
19 | 13.3 | 3.2 | 7.2 | 5.4

Our models are able to reproduce the magnitude of the peak luminosity,
although exact matching of the lightcurve post-peak is difficult. Phases of
the swept up wind becoming transparent, followed by the ejecta can be seen as
the sudden drop-offs in Fig. 15. We find that the width of the Event B peak is
dependent to some degree to how the density of the wind varies with distance
from the star. The figure shows the resultant models where we assume
$\rho_{\rm wind}(r)\propto r^{-1.6}$. We found that the shallower the density
gradient the wider the peak and a best match is found with an exponent
$n=-1.6$. In general, the models that best match the supernova lightcurve have
low ejecta masses on the order of 1 to 2 M⊙. Some models which have
experienced a merger during their binary evolution have a higher ejecta mass
do not match the lightcurve, being less luminous or evolving more slowly.
Achieving an exact match between the models and observed lightcurve would
require significant fine-tuning of the details of the CSM around the star, in
terms of density profile, wind velocity, and the details of the wind
acceleration. An exact match may also be impossible given the spherically
symmetric assumptions of SNEC. However, we take the reasonably close match
between the model and observed lightcurves to indicate that a subset of the
BPASS models can explain AT 2016jbu.

Intriguingly, the low $CO$ core mass of several of the progenitor models
suggest an explosion close to the electron-capture regime where lower nickel
masses and explosion energies would be expected.

###  7.4 Was AT 2016jbu a Core-Collapse Supernova?

The main point of controversy is whether AT 2016jbu and SN 2009ip-like
transients are indeed CCSNe; meaning the progenitor has been destroyed and the
transient will eventually decay following a radioactive decay tail. This begs
the question; If these are indeed CCSNe, when did core-collapse occur?

SN 2009ip-like transients display two broad, luminous events, rather than the
singularly peaked lightcurve, typically associated with SNe. Mauerhan et al.
(2013b) suggest that Event A is a CCSN and Event B is a result of ejecta
interacting with a dense CSM. In this scenario, with respect to AT 2016jbu,
the duration of Event A ($\sim$ 60 days) is the time needed for this ejecta to
reach the inner edge of the CSM. This scenario would be consistent with the
early evolution of $R_{\rm BB}$ expanding at $\sim$700 km s-1; however this
velocity is implausibly slow for SN ejecta. More problematic still, at $-21$ d
we see an increase in velocity where $R_{\rm BB}$ expands at $\sim$4500 km
s-1. In the case of core-collapse, we hence regard it as more plausible that
Event B is the terminal explosion of the progenitor, where the ejecta
interacts with a non-terminal outburst that was ejected at $\sim 700$ km s-1
around the start of Event A. This scenario is also reinforced by the rise time
($\sim 17$ days) and peak magnitude ($M_{V}\sim-18.5$ mag) of Event B (Nyholm
et al., 2020).

We find a low value of 56Ni of $\lesssim$ 0.016 M⊙ for AT 2016jbu, consistent
with other SN 2009ip-like transients. Such a low 56Ni mass would be unusual
for a normal CCSN, although an exception would be a faint electron capture SN
(ECSN) or a sub-luminous Fe CCSN from a star with a ZAMS mass of around 8 – 10
M⊙. However, we find the mass of the AT 2016jbu progenitor to be significantly
larger than that expected for an ECSN progenitor (Doherty et al., 2017).
Additionally the inferred explosion energy of $5.6\times 10^{50}$ ergs (which
may be a lower limit, as spherical symmetry is assumed) is too high for a
typical ECSN Wanajo et al. (2009). A final possibility that can explain such a
low Ni mass (if this is a CCSN) is significant fallback onto a compact remnant
(Zampieri et al., 1998; Benetti et al., 2016).

Some challenges remain for the fallback scenario. A low metallicity
environment is required, so that the progenitor star has retained much of its
ZAMS mass (e.g. Heger et al., 2003). This is hence an appealing scenario for
SN 2009ip, due to its remote location ($\sim$5 kpc from its host Smith et al.,
2016b) and naturally low-metallicity environment. Conversely, this contradicts
what we see for the environment around AT 2016jbu in Sect. 4.2, where we find
an approximate solar metallicity of 8.66 dex. It is hence expected that a
$\sim$20 M⊙ progenitor will loose a significant fraction of its mass before
exploding.

We see from Fig. 9 that AT 2016jbu is located near a moderately star-forming
region that is likely to host CCSNe, as seen from SN 1999ga. On the contrary,
SN 2009ip is located on the outskirts of its host spiral galaxy, NGC 7259, at
a galactocentric radius of $\sim$ 5 kpc. Smith et al. (2016b) finds no strong
indication of massive star formation anywhere in the vicinity around SN
2009ip, unlike what is seen for AT 2016jbu. If the progenitor of SN 2009ip and
AT 2016jbu are similar, as is suggested by their photometric and spectral
evolution, then this begs the question why SN 2009ip is on its own.

One of the biggest difficulties with AT 2016jbu as a CCSN is that it is in
stark contrast to the predictions of single star stellar evolutionary models.
A 20 M⊙ star is expected to end its life as a RSG which undergoes Fe core-
collapse (Heger et al., 2003). From our dusty modelling in Sect. 2, we find
that the progenitor of AT 2016jbu is not situated at the end of any single
star evolutionary track. This suggests that the progenitor is not sufficiently
evolved to undergo core-collapse. Our conclusion in Sect. 2 also suggests that
the progenitor of AT 2016jbu is not a RSG but rather a YHG. We also note that
if AT 2016jbu is indeed a CCSN, it is more appropriate to compare to the
luminosity of the progenitor to the terminal luminosity of the models
(typically corresponding to the end of core He-burning), in which case we find
that it must have been a 12–16 M⊙ star. One must caution however that if the
progenitor of AT 2016jbu was in a binary, then the expectations from single
star evolution can be drastically altered. However, even if AT 2016jbu does
arise from a binary progenitor system, models do not necessarily predict
outbursts or eruptions immediately prior to explosion as seen in this case
(discussed further in Sect. 7.5 below). Clearly, further detailed stellar
evolutionary modeling is required to fully explain the progenitor (or
progenitor system) of AT 2016jbu.

A tantalising hint of a surviving progenitor is AT 2016jbu returning to its
pre-explosion magnitude in 2019, as shown in Fig. 2. However, this detection
may be serendipitous and further late time monitoring will be needed to
confirm any surviving progenitor.

###  7.5 Binary Interaction

Several authors have suggested that SN 2009ip-like transients are a result of
binary interaction (Smith et al., 2014b; Kashi et al., 2013; Soker & Kashi,
2013; Smith et al., 2018) as well as some other SN Impostors e.g. SN 2000ch
(Pastorello et al., 2010; Smith, 2011; Clark et al., 2013). Mass transfer
within a binary system could naturally explain an asymmetric CSM environment,
which we interpret as a circumstellar/circumbinary disk for AT 2016jbu.

Smith & Tombleson (2015) suggest that the isolated location of SN 2009ip may
be explained as they are Kicked Mass Gainers in a binary star system. For the
progenitor to travel $\sim$ 5 kpc within the required lifespan of a 50 – 80 M⊙
star, a binary companion may be required to provide an additional source of
fuel after the stars have been ejected (Smith et al., 2016b).

Binary merger events have recently been associated with Red Novae (RNe) and
the more extreme, Luminous Red Novae (LRNe) (see review by Pastorello et al.,
2019b, and references therein). These transients typically fall into the class
of Gap Transients (Kasliwal et al., 2012; Pastorello & Fraser, 2019) and are
amongst the most powerful stellar cataclysms. LRNe span a wide range of
absolute magnitudes, from $-4$–$-15$ mag (Pastorello & Fraser, 2019), and show
a wide range of lightcurve shape and duration.

The physical interpretation of LRNe is debated. The progenitors of LRNe are
likely massive contact binaries, and the doubled peaked lightcurve is a
consequence of a stellar merger plus a Common Envelope Ejection (CEE) (Smith
et al., 2016a; Metzger & Pejcha, 2017; Pastorello et al., 2019b). Pastorello
et al. (2019b) suggest that there may be a continuum spanning between RNe to
LRNe, with the possibility that this range can reach to brighter magnitudes
(most likely caused by higher mass systems). SN 2009ip-like events may be some
combination of a binary merger where the system consists of a relatively
massive primary where the stars undergo a Common Envelope (CE) phase followed
by a massive eruption.

AT 2016jbu and SN 2009ip-like transients show a common peak magnitude and
shape (i.e. Event B appear to be similar among SN 2009ip-like events). We do
not have adequate colour information for the peak of Event A for AT 2016jbu
however, Event B has a colour of B-V$\sim$ 0, which is comparable to that seen
in LRNe in their first peak. AT 2016jbu never gets redder than $B-V$ of
$\lesssim$ 0.8 and after $\sim$ 1.2 years, when the transient returns to a
$B-V$ value of $\sim$ 0.2. If AT 2016jbu is indeed related to LRNe, continued
interaction in AT 2016jbu may be responsible for the relatively blue colour at
late times.

Soker & Kashi (2013) proposed that SN 2009ip is the result of the merger of a
massive LBV with a binary companion in their “mergerburst” model. This model
agrees with observations of SN 2009ip quite well, such as the moderate ejecta
mass (a few M⊙), most of which is moving at less than 5000 km s-1. They
further predict the remnant of their mergerburst will be a hot red giant star
that will become apparent years after the transient fades, as is commonly
associated with RNe and LRNe (e.g Pastorello et al., 2019b). Kashi et al.
(2013) discuss a similar explosion mechanism to the scenario we discuss in
Sect. 5.1 and conclude the double-peaked lightcurve of SN 2009ip may be
explained by to two successive outburst, separated by $\sim 20$ days caused by
periastron passages of the binary system.

It is appealing to conclude that AT 2016jbu is the result of a coalescing
binary. This can naturally explain the historic variability, double-peaked
lightcurve, and (inferred) asymmetric CS environment i.e. disk or bipolar
outflow. Metzger & Pejcha (2017) proposed that LRNe can be well modeled by a
single symmetric eruption in an asymmetric CSM environment. This asymmetric
CSM is fueled by mass transfer within the binary over many orbits preceding
the double-peaked event. The first peak of LRNe can be comfortably powered via
cooling envelope emission from fast moving ejecta. Radiative shocks from the
collision of this ejecta with material in the equatorial plane then power the
second peak. This would be inconsistent with our proposed “catch-up” scenario
for AT 2016jbu, although it cannot be ruled out conclusively.

We can speculate that the events prior to Event B in AT 2016jbu and SN 2009ip-
like events are similar to LRNe, including as mass transfer / Roche Lobe
Overflow (RLOF) seen in the decade leading up to Event A, and a merger/CEE
powering Event A itself. To explain the homogeneity of Event B, the merging of
the binary system must cause a violent (and possibly terminal) eruption.

Each SN 2009ip-like transient remains relatively blue for a long period of
time, unlike what we see in LRNe, which is likely a sign of continued
interaction. If we assume that SN 2009ip-like transients are indeed an
upscaled version of LRNe, then this continued interaction at late times may
reflect a more massive progenitor than is commonly associated with LRNe. In
this scenario would expect a surviving star to become visible after this
interaction has abated.

##  8 Conclusion

In this paper, we have investigated the progenitor and environment of AT
2016jbu as well as modelling the transient itself. If AT 2016jbu is a single
star, we find that the progenitor is consistent with a $\sim$ 22 M⊙ progenitor
(e.g. Fig. 4 in Smartt et al., 2009), with a color consistent with a YHG,
roughly consistent with K18. Modelling of circumstellar dust using dusty gives
a luminosity and temperature of the progenitor similar to known YHGs. We show
that the local environment around the progenitor of AT 2016jbu is consistent
with a CCSN from a progenitor with ZAMS mass $\sim$ 20 M⊙; as the stellar
population has an age of 15–200 Myr. We confidently rule out the possibility
that the progenitor of AT 2016jbu is an LBV of 50–80 M⊙, as has been proposed
for SN 2009ip (Smith et al., 2010; Foley et al., 2011).

We find that the Event A/B light curve can be modelled by two shells of
material, with the later Event B being powered by a “catch-up” scenario,
involving two eruptive mass loss events and pre-existing CSM. Spectroscopic
and photometric evolution is consistent with a spherically symmetric ejecta
colliding with, and temporarily engulfing, previously ejected, asymmetric
material. This interaction is the dominant energy source after $\sim$ 2
months. After $\sim$ 200 days, AT 2016jbu shows increased interaction, likely
reflecting a clumpy CSM.

AT 2016jbu shows tentative evidence for core-collapse. We find a upper limit
of 56Ni of $\lesssim 0.016$ M⊙ but with strong on-going CSM interaction at
this time, the real value of 56Ni is probably much lower (if any at all).
Almost 1.5 years after maximum brightness, AT 2016jbu lacks signs of
explosively nucleosynthesised material or emission from the burning products
of late time stellar evolution.

We explore the possibility that AT 2016jbu is the result of a binary system.
We compare our progenitor models with an extensive group of BPASS models,
exploring both CCSN and non-terminal events. We find that matching models have
MZAMS $\lesssim~{}26$ M⊙. Steady state mass loss due to the progenitor wind is
unable to produce the CSM density necessary to power the lightcurve and
episodic mass loss may be required. Using SNEC we find that a relatively low
explosion energy
($5.5\text{\times}{10}^{49}\text{\,}\mathrm{e}\mathrm{r}\mathrm{g}\mathrm{s}$)
with a small ejecta mass ($\sim$ 1-2 M⊙) can comfortably power AT 2016jbu
(assuming spherical symmetry). If we account for a high degree of asymmetry,
we may have an explosion energy on par with a typical CCSN.

It appears that there is not a simple explanation for these transients.
Following Hickam’s dictum, a low energy SN within a binary system with a disk-
like CSM can account for the rise and peak of Event B, low 56Ni, continued CSM
interaction, and unique spectral features of AT 2016jbu. Additional binary
interaction might explain Event A e.g. due to a merger or CEE. Detailed
modelling of this proposed scenario is beyond the scope of this paper and
future work will involve exploring these scenarios in a non-symmetric setting.

The true nature of AT 2016jbu (and other SN 2009ip-like transients) remains
elusive. Perhaps the ultimate answer will come if or when very late time
observations reveal a surviving progenitor. To date, no conclusive evidence
exists as to whether these transients destroy their progenitor. However, one
must account for the possibility that if the progenitor survived, it may be
obscured by a significant amount of dust. Deep images covering the full SED
will hence be required to confidently rule out surviving, but dust-enshrouded,
star. To this end, future observations with the upcoming James Webb Space
Telescope will be essential. Alongside this, deep optical imaging from the
Vera C. Rubin Observatory may capture similar pre-explosion variability in the
years/decades prior to future SN 2009ip-like events, perhaps even allowing for
a countdown timer before these events.

## Acknowledgements

S. J. Brennan acknowledges support from Science Foundation Ireland and the
Royal Society (RS-EA/3471). M.F is supported by a Royal Society - Science
Foundation Ireland University Research Fellowship. T.M.B was funded by the
CONICYT PFCHA / DOCTORADOBECAS CHILE/2017-72180113 and acknowledges their
financial support from the Spanish Ministerio de Ciencia e Innovación (MCIN),
the Agencia Estatal de Investigación (AEI) 10.13039/501100011033 under the
PID2020-115253GA-I00 HOSTFLOWS project, and from Centro Superior de
Investigaciones Científicas (CSIC) under the PIE project 20215AT016. K.M. is
funded by the EU H2020 ERC grant no. 758638. T.W.C acknowledges the EU Funding
under Marie Skłodowska-Curie grant H2020-MSCA-IF-2018-842471, and thanks to
Thomas Krühler for GROND data reduction. M.N is supported by a Royal
Astronomical Society Research Fellowship. B.J.S is supported by NSF grants
AST-1908952, AST-1920392, AST-1911074, and NASA award 80NSSC19K1717. M.S is
supported by generous grants from Villum FONDEN (13261,28021) and by a project
grant (8021-00170B) from the Independent Research Fund Denmark. L.H
acknowledges support for Watcher from Science Foundation Ireland grant
07/RFP/PHYF295. Time domain research by D.J.S. is supported by NSF grants
AST-1821987, 1813466, & 1908972, and by the Heising-Simons Foundation under
grant #2020-1864. N.E.R. acknowledges support from MIUR, PRIN 2017 (grant
20179ZF5KS). Support for JLP is provided in part by ANID through the Fondecyt
regular grant 1191038 and through the Millennium Science Initiative grant
ICN12$\\_$009, awarded to The Millennium Institute of Astrophysics, MAS. L.G.
acknowledges financial support from the Spanish Ministry of Science,
Innovation and Universities (MICIU) under the 2019 Ramón y Cajal program
RYC2019-027683 and from the Spanish MICIU project PID2020-115253GA-I00. D.A.H
and D.H are supported by AST-1911151, AST19-11225, and NASA Swift grant
80NSSC19K1639. G.P acknowledge support by the Ministry of Economy,
Development, and Tourism’s Millennium Science Initiative through grant
IC120009, awarded to The Millennium Institute of Astrophysics, MAS. L.T.
acknowledges support from MIUR (PRIN 2017 grant 20179ZF5KS). Support for TW-SH
was provided by NASA through the NASA Hubble Fellowship grant HST-
HF2-51458.001-A awarded by the Space Telescope Science Institute, which is
operated by the Association of Universities for Research in Astronomy, Inc.,
for NASA, under contract NAS5-26555. H.K. was funded by the Academy of Finland
projects 324504 and 328898.

This research made use of Astropy888http://www.astropy.org, a community-
developed core Python package for Astronomy (Astropy Collaboration et al.,
2013; Price-Whelan et al., 2018). This research made use of data provided by
Astrometry.net999https://astrometry.net/use.html. Parts of this research were
supported by the Australian Research Council Centre of Excellence for All Sky
Astrophysics in 3 Dimensions (ASTRO 3D), through project number CE170100013.
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. We acknowledge Telescope Access Program (TAP) funded by the
NAOC, CAS, and the Special Fund for Astronomy from the Ministry of Finance.
This work was partially supported from Polish NCN grants: Harmonia No.
2018/30/M/ST9/00311 and Daina No. 2017/27/L/ST9/03221. This work made use of
v2.2.1 of the Binary Population and Spectral Synthesis (BPASS) models as
described in Eldridge et al. (2017) and Stanway & Eldridge (2018). This
research is based on observations made with the NASA/ESA Hubble Space
Telescope obtained from the Space Telescope Science Institute, which is
operated by the Association of Universities for Research in Astronomy, Inc.,
under NASA contract NAS 5-26555. These observations are associated with
program 15645. Observations were also obtained from the Hubble Legacy Archive,
which is a collaboration between the Space Telescope Science Institute
(STScI/NASA), the Space Telescope European Coordinating Facility (ST-
ECF/ESAC/ESA) and the Canadian Astronomy Data Centre (CADC/NRC/CSA). This
research has made use of the SVO Filter Profile
Service101010http://svo2.cab.inta-csic.es/theory/fps/ supported from the
Spanish MINECO through grant AYA2017-84089.

##  Appendix A Author Affiliations

1 School of Physics, O’Brien Centre for Science North, University College
Dublin, Belfield, Dublin 4, Ireland
2 The Oskar Klein Centre, Department of Physics, AlbaNova, Stockholm
University, SE-106 91 Stockholm, Sweden
3 INAF-Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122
Padova, Italy
4 Department of Physics and Astronomy, University of Turku, FI-20014, Turku,
Finland
5 The Department of Physics, The University of Auckland, Private Bag 92019,
Auckland, New Zealand
6 The Oskar Klein Centre, Department of Astronomy, AlbaNova, Stockholm
University, SE-106 91 Stockholm, Sweden
7 Max-Planck-Institut für Extraterrestrische Physik, Giessenbachstraße 1,
85748 Garching, Germany
8 Department of Astronomy, The Ohio State University, 140 W. 18th Avenue,
Columbus, OH 43210, USA
9 Center for Cosmology and AstroParticle Physics (CCAPP), The Ohio State
University, 191 W. Woodruff Avenue, Columbus, OH 43210, USA
10 Department of Physics and Astronomy, Texas A$\&$M University, 4242 TAMU,
College Station, TX 77843, USA
11 Cerro Tololo Inter-American Observatory, NSF’s National Optical-Infrared
Astronomy Research Laboratory, Casilla 603, La Serena, Chile
12 Institut d’Astrophysique de Paris (IAP), CNRS $\&$ Sorbonne Universite,
France
13 Kavli Institute for Astronomy and Astrophysics, Peking University, Yi He
Yuan Road 5, Hai Dian District, Beijing 100871, China
14 NINAF-Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5,
I-35122 Padova, Italy
15 Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans
s/n, 08193 Barcelona, Spain
16 Steward Observatory, University of Arizona, 933 North Cherry Avenue,
Tucson, AZ 85721-0065, USA
17 Department of Physics, Florida State University, 77 Chieftan Way,
Tallahassee, FL 32306, USA
18 Tuorla Observatory, Department of Physics and Astronomy, FI-20014
University of Turku, Finland.
19 Finnish Centre for Astronomy with ESO (FINCA), FI-20014 University of
Turku, Finland
20 CBA Kleinkaroo, Calitzdorp, South Africa
21 Departamento de Ciencias Físicas, Universidad Andres Bello, Avda. Republica
252, Santiago, 8320000, Chile
22 Millennium Institute of Astrophysics, Santiago, Chile
23 Department of Astronomy/Steward Observatory, 933 North Cherry Avenue, Rm.
N204, Tucson, AZ 85721-0065, USA
24 Institute for Astronomy, University of Hawai’i, 2680 Woodlawn Drive,
Honolulu, HI 96822, USA
25 Astrophysics Research Centre, School of Maths and Physics, Queen’s
University Belfast, Belfast BT7 1NN, UK
26 Mt Stromlo Observatory, The Research School of Astronomy and Astrophysics,
Australian National University, ACT 2601, Australia
27 National Centre for the Public Awareness of Science, Australian National
University, Canberra, ACT 2611, Australia
28 The ARC Centre of Excellence for All-Sky Astrophysics in 3 Dimension (ASTRO
3D), Australia
29 Astronomical Observatory, University of Warsaw, Al. Ujazdowskie 4, 00-478
Warszawa, Poland
30 Institute of Astronomy, Madingley Road, Cambridge,CB3 0HA, UK
31 Unidad Mixta Internacional Franco-Chilena de Astronomía, CNRS/INSU UMI 3386
and Instituto de Astrofísica, Pontificia Universidad Católica de Chile,
Santiago, Chile
32 Aix Marseille Univ, CNRS, CNES, LAM, Marseille, France
33 RHEA Group for ESA, European Space Astronomy Centre (ESAC-ESA), Madrid,
Spain
34 nstitute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans,
s/n, E-08193 Barcelona, Spain.
35 Tuorla Observatory, Department of Physics and Astronomy, FI-20014
University of Turku, Finland
36 Kavli Institute for Cosmology, Institute of Astronomy, Madingley Road,
Cambridge, CB3 0HA, UK
37 Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA
93117-5575, USA
38 Department of Physics, University of California, Santa Barbara, CA
93106-9530, USA
39 The Observatories of the Carnegie Institution for Science, 813 Santa
Barbara St., Pasadena, CA 91101, USA
40 School of Physics $\&$ Astronomy, Cardiff University, Queens Buildings, The
Parade, Cardiff, CF24 3AA, UK
41 School of Physics and Astronomy, University of Southampton, Southampton,
Hampshire, SO17 1BJ, UK
42 School of Physics, Trinity College Dublin, The University of Dublin, Dublin
2, Ireland
43 Department of Physics, University of the Free State, PO Box 339,
Bloemfontein 9300, South Africa
44 Carnegie Observatories, Las Campanas Observatory, Colina El Pino, Casilla
601, Chile
45 Birmingham Institute for Gravitational Wave Astronomy and School of Physics
and Astronomy, University of Birmingham, Birmingham B15 2TT, UK
46 Institute for Astronomy, University of Edinburgh, Royal Observatory,
Blackford Hill, EH9 3HJ, UK
47 Nucleo de Astronomıa de la Facultad de Ingenierıa y Ciencias, Universidad
Diego Portales, Av. Ej´ercito 441, Santiago, Chile
48 Department of Physics and Astronomy University of North Carolina at Chapel
Hill Chapel Hill, NC 27599, USA
49 Department of Physics, Florida State University, 77 Chieftain Way,
Tallahassee, FL 32306-4350, USA
50 Department of Physics and Astronomy, Aarhus University, Ny Munkegade,
DK-8000 Aarhus C, Denmark
51 Department of Physics, University of Warwick, Coventry CV4 7 AL, UK
52 Astrophysics Research Centre, School of Mathematics and Physics, Queen‘s
University Belfast, Belfast BT7 1NN, UK
53 Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans, s/n, E-08193 Barcelona, Spain. 54 Center for Astrophysics | Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138-1516, USA 55 The NSF AI Institute for Artificial Intelligence and Fundamental Interactions

## Data availability

The photometric and spectroscopic data underlying this article are available
as described in Paper I. The BPASS models are available at
https://bpass.auckland.ac.nz/, while HST data are available at the Mikulski
Archive for Space Telescopes athttps://archive.stsci.edu.

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