# Assessing the impact of valence sd neutrons and protons on fusion

Varinderjit Singh  J. Vadas  T.K. Steinbach  B. B. Wiggins  S. Hudan  R. T.
deSouza111Corresponding author Department of Chemistry and Center for
Exploration of Energy and Matter, Indiana University
2401 Milo B. Sampson Lane, Bloomington, IN 47408, USA

###### Abstract

Experimental near-barrier fusion cross-sections for 17F + 12C are compared to
the fusion excitation functions for 16,18O, 19F, and 20Ne ions on a carbon
target. Comparison of the reduced fusion cross-section for the different
systems accounts for the differing static size of the incident ions and
changes in fusion barrier. Remaining trends of the fusion cross-section above
the barrier are observed. These trends are interpreted as the interplay of the
sd protons and neutrons. The experimental data are also compared to a widely-
used analytic model of near-barrier fusion, a time-dependent Hartree-Fock
model, and coupled channels calculations.

###### keywords:

near-barrier fusion , RIB fusion , fusion enhancement

††journal: Physics Letters B

##  1 Introduction

Nuclear fusion is a topic of considerable interest both from a fundamental
perspective as well as in the field of nuclear astrophysics [1]. Nuclei just
beyond a closed shell present a unique opportunity to probe the interplay of
shell and collective effects on the fusion process. In particular, light
nuclei just beyond the 1p1/2 shell, namely isotopes of oxygen, fluorine, and
neon are good candidates for examination. In this work, the fusion of various
isotopes of these elements with a carbon target at near-barrier energies is
examined. The results of this work, which combines both stable and radioactive
beams, points to the potential of low-energy beams at radioactive beam
facilities [2, 3] for examining the impact of neutron-excess on fusion.

Addition of neutrons and protons just beyond the 1p1/2 shells of 16O clearly
changes both the matter and charge distributions of the nuclei. Theoretical
calculations indicate that for a large neutron excess, e.g. 24O as compared to
16O, fusion with 16O target is significantly enhanced [4]. The impact of
adding just a few neutrons or protons beyond the 1p${}_{\frac{1}{2}}$ shell on
fusion is less clear. With increased atomic or mass number, the fusion barrier
and consequently the fusion cross-section is clearly impacted. We propose to
go beyond these trivial systematic differences and examine the detailed
differences in the fusion cross-section. A particularly interesting case to
investigate is fusion of the nucleus 17F which exhibits a proton-halo when in
its 2s1/2 excited state [5]. It is presently unclear whether an increased
radial extent results in a fusion enhancement or weak binding results in a
decreased fusion cross-section. For the case of 17F + 208Pb, neither an
enhancement nor a suppression of fusion was observed relative to 19F + 208Pb
[7]. However, in the case of fusion with a large target nucleus, such as
208Pb, the impact of adding two neutrons on fusion might be diminished. To
examine the impact on fusion of adding a few protons and neutrons to the sd
shell we measure fusion for 17F + 12C and compare it with fusion induced by O,
F, and Ne beams on 12C.

##  2 Experimental Details

Discovery of halo nuclei, notably 11Li [8, 9] was achieved through systematic
examination of the interaction cross-sections for lithium isotopes. At high
incident energy one expects the sudden approximation to be valid in describing
the collision of the nuclei. Consequently, the nuclear densities do not have
enough time to rearrange as the projectile and target nuclei come into
contact. Thus a measurement of the interaction cross-section at high energy
probes the nuclear size and other geometrical features such as deformation,
all of which are considered “static” [8]. Hence, the measured interaction
cross-section, $\sigma_{I}$, provides a direct and effective measure of the
extent of the matter distribution.

To better understand the change in the static size of the different nuclei
considered in this work, we examine the interaction cross-sections measured at
high energy. Presented in Table 1 are $\sigma_{I}$ for O, F, and Ne nuclei
with a carbon target [10] with the number of protons and neutrons in the sd
shell indicated. The closure of the 1p${}_{\frac{1}{2}}$ with N=8 provides a
natural reference from which to examine the impact made by the presence of a
few nucleons in the sd shell. Addition of two neutrons to 16O, 17F, and 18Ne
increases the interaction cross-section by 50, 61, and 68 mb respectively.
Addition of two protons to 16O, i.e. 18Ne, results in an increase of
$\sigma_{I}$ by 94 mb while in the case 18O to 20Ne, it increases by 112 mb.
The difference in the interaction cross-section with the addition of two
protons (94-112 mb) as compared to addition of two neutrons (50-68 mb) is
presumably due to the repulsion of the two protons with the 16O core. Such a
difference is not captured by systematics which relate the radius of a nucleus
to simply the mass number, A, suggesting the advantage of using $\sigma_{I}$
to account for trivial size effects.

Table 1: Interaction cross-sections for oxygen, fluorine, and neon nuclei with a carbon target at an energy about 900 A MeV. Taken from [10]. 18Ne2p0n |  20Ne2p2n
---|---
$\sigma_{I}$=1076$\pm$25 mb |  $\sigma_{I}$=1144$\pm$10 mb
17F1p0n |  19F1p2n
$\sigma_{I}$=982$\pm$32 mb |  $\sigma_{I}$=1043$\pm$24 mb
16O0p0n |  18O0p2n
$\sigma_{I}$=982$\pm$6 mb |  $\sigma_{I}$=1032$\pm$26 mb

While the excitation function for fusion of 16,18O, 19F, and 20Ne ions with a
carbon target already exists, no data exists for 17F+12C or 18Ne+12C. A beam
of 16O ions at Florida State University was used to produce 17F ions via a
(d,n) reaction. These ions were separated from the incident beam by the
electromagnetic spectrometer RESOLUT [11]. The beam exiting the spectrometer
consisted of both 17F and residual 16O ions necessitating identification of
each ion on a particle-by-particle basis. The presence of both species allowed
the simultaneous measurement of 16O + 12C and 17F + 12C providing a built-in
reference measurement and comparison with the well-established fusion
excitation function for 16O + 12C.

Figure 1: Schematic of the experimental setup along with a PID and ETOF
spectra. Particle identification of the 17F and 16O ions incident on the
target is displayed in the left plot along with a representative ETOF spectrum
(selected on 17F) used to identify evaporation residues (ER) in the right
plot.  Figure 2: Fusion excitation functions for 16O + 12C (top panel) and
17,19F + 12C (bottom panel). The predictions of an analytic fusion model [12]
are indicated by the solid lines.

The setup used to measure fusion of fluorine and oxygen ions with carbon
nuclei in this experiment is depicted in Fig. 1. Two microchannel plate
detectors designated MCPRESOLUT and MCPTGT spaced $\sim$3.5 m apart and a
compact ionization detector (CID) permitted a $\Delta$E-TOF measurement for
each ion incident on the target. Measurement of the $\Delta$E-TOF (Fig. 1)
clearly shows three peaks corresponding to 17F9+, 16O7+, and 16O8+ ions. The
intensity of the 17F beam incident on the target was 3-7 x 103 ions/s with a
purity of 37%-54%.

Fusion of a 17F (16O) nucleus in the beam together with a 12C nucleus in the
target foil results in the production of an excited 29P (28Si) nucleus. For
near-barrier collisions the excitation of the fusion product is relatively
modest, E∗ $\approx$ 30 MeV. Emission of particles as the fusion product de-
excites, deflects the evaporation residue (ER) from the beam direction. This
deflection allows its detection and identification using two annular silicon
detectors designated T2 and T3 that subtend the angular range 3.5∘ $<$
$\theta_{lab}$ $<$ 25∘. Using the measured energy deposit in the silicon
detectors and the time-of-flight [13], the mass of the ion was calculated. ERs
were cleanly distinguished from beam particles by their mass[14, 15]. A
representative ETOF spectrum is presented in Fig. 1.

The fusion cross-section, $\sigma_{F}$ is extracted from the measured yield of
evaporation residues through the relation $\sigma_{F}$ = NER/($\epsilon_{ER}$
x t x NI) where NI is the number of beam particles of a given type incident on
the target, t is the target thickness, $\epsilon_{ER}$ is the detection
efficiency, and NER is the number of evaporation residues detected. The number
NI is determined by counting the particles with the appropriate time-of-flight
between the two microchannel plates that additionally have the correct
identification in the $\Delta$E-TOF map depicted in Fig. 1. The target
thickness, t, is 105 $\pm$ 0.5 $\mu$g/cm2. The number of detected residues,
NER, is determined by summing the number of detected residues clearly
identified by the ETOF technique as shown in Fig. 1 [15]. To obtain the
detection efficiency, $\epsilon_{ER}$, a statistical model is used to describe
the de-excitation of the fusion product together with the geometric acceptance
of the experimental setup. The detection efficiency varied from $\sim$81% at
the highest incident energies measured to $\sim$85% at the lowest incident
energy due to the changing kinematics of the reaction with an associated
uncertainty of approximately 5$\%$.

Figure 3: Fusion excitation functions for 18O+12C (top panel) and 20Ne+12C
(bottom panel) together with the predictions of the analytic model[12]. The
fit of the 20Ne data by a one-dimensional barrier-penetration model [16] is
depicted by the dashed line.

##  3 Results and Discussion

Depicted in Fig. 2 are the measured excitation functions for fusion of 16O and
17F ions with a carbon target. With the experimental technique used, both the
16O + 12C and 17F + 12C excitation functions are simultaneously measured. In
the top panel, for 16O, one observes the good agreement of this measurement,
depicted as the solid (red) circles, with the experimental cross-sections
reported in the literature [17, 18, 19, 20]. The present measurement
reproduces the known resonances in the interval 10 MeV $\leq$Ec.m.$\leq$ 18
MeV demonstrating that the measurement of the fusion excitation function for
17F + 12C is robust. In the lower panel of Fig. 2 the measured cross-sections
for 17F+12C (solid red circles) are presented along with the excitation
function for 19F+12C. The multiple measurements for 19F [21, 19, 22] are in
agreement within the measurement uncertainties. As might be expected naively
by the reduction of two neutrons, the 17F-induced fusion of this measurement
exhibits a lower cross-section than 19F for all energies shown. The measured
excitation functions are compared with the predictions of an analytic model
based on a parameterization of the Sao Paulo potential model coupled with a
barrier penetration formalism [12, 23, 24]. This model which has parameterized
a large number of reactions for low and mid-mass systems is a useful tool for
network simulations in the near-barrier regime. Despite having several fitted
parameters, the analytic model has no adjustable parameters and thus serves as
a useful benchmark of the expected systematic behavior. The above-barrier
cross-sections predicted by this model are depicted in Fig. 2 as the solid
lines. In the case of the 16O fusion the analytic model provides a reasonable
description of the overall behavior of the excitation function. The failure of
the analytic model to reproduce the resonance is unsurprising as it assumes a
smooth, structureless barrier. For the fluorine isotopes, while the excitation
function for 17F is reasonably described, the excitation function for 19F is
clearly underpredicted in the energy range measured. Nonetheless, the analytic
model does predict an increase in the cross-section associated with the
presence of the two additional neutrons in 19F as compared to 17F.

Also presented in Fig. 2 is the excitation function for 17F + 12C recently
measured by Asher et al. (solid blue triangles) [25]. The statistical
uncertainties of the Asher measurement are significantly larger than those of
the present measurement. As importantly, the Asher measurement has
significantly larger uncertainty in the energy, an inherent consequence of the
technique utilized. Even considering the uncertainties of both excitation
functions it is clear that the Asher excitation function reports larger cross-
sections or is shifted to lower energies as compared to the present
measurement. Surprisingly, the excitation function measured by Asher for
Ec.m.$\leq$15 MeV is essentially consistent with the excitation function for
19F. This result would imply that the removal of two neutrons from 19F makes a
negligible difference in both the barrier and the size of the fusing system.
Such a result is inconsistent with the analytic model results which describes
the expected systematic behavior. Consequently, for the remainder of this work
only the present 17F data is considered.

The excitation functions for fusion of 18O and 20Ne nuclei with carbon are
shown in Fig. 3. While good agreement is observed for the various 18O
datasets, in the case of 20Ne, for 12 MeV $\leq$Ec.m.$\leq$15 MeV a
discrepancy exists between the measurements of Hulke [26] and deYoung [27].
This discrepancy at Ec.m. $\sim$15 MeV is larger than the reported
uncertainties by both experimental measurements. Moreover, the shape of the
excitation function determined by Hulke [26] deviates from the behavior of a
smooth barrier as indicated by the analytic model [12]. The Hulke data [26]
for Ec.m.$\leq$11 MeV and the deYoung data [27] can be described by a one-
dimensional barrier penetration model [16] as indicated by the dashed line in
Fig. 3.

In order to appropriately compare all the fusion excitation functions for
these light nuclei, we scale the fusion cross-section by the interaction
cross-section, $\sigma_{I}$, presented in Table 1. In addition, the trivial
effect of the different barriers is accounted for by examining the reduced
cross-section, $\sigma_{F}$/$\sigma_{I}$, as a function of the above-barrier
energy, Ec.m.-VB. The value of the barrier, VB, is taken from the Bass model
[28]. The uncertainties shown in the reduced cross-section reflect the
uncertainties in both $\sigma_{F}$ and $\sigma_{I}$. This presentation allows
one to investigate differences between the nuclei shown after effectively
eliminating systematic differences in the static size and barrier. Comparison
of all the excitation functions in Fig. 4 yields some interesting results. In
Fig. 4a one observes that for (Ec.m.-VB)$\leq$9 MeV, the reduced cross-
sections for 16O and 18O are relatively comparable. The prominent resonance
structure for 16O-induced fusion is absent in the case of 18O. Above 9 MeV one
observes that the reduced cross-section for 18O exceeds that of 16O. This
increase in the reduced cross-section indicates the impact of the two sd
neutrons on the fusion cross-section over and above the increase in the static
size. In further understanding the role of sd valence nucleons on fusion, we
elect to choose 18O as our reference. The absence of strong, sharp resonance
structures in the 18O reaction supports this choice. To facilitate the use of
18O as a reference we have described the data by the smooth curve depicted in
Fig. 4a. This curve corresponds to a third-order polynomial fit to the reduced
cross-sections for 18O and simply serves as an adequate representation of the
18O data in Fig. 4b.

Figure 4: Comparison of the reduced excitation functions for fusion of 16O and
18O ions (top panel) and 18O, 20Ne, 19F, and 17F ions (lower panels) on a
carbon target. The line through the 20Ne data simply serves to guide the eye.

Presented in Fig. 4b are the reduced excitation functions for 19F, 20Ne, and
17F in comparison to 18O. Within the measurement uncertainties the 19F data
manifests the same reduced cross-section as the 18O. However, there is a
subtle indication that the reduced cross-section for 19F appears to be
systematically slightly lower than that of 18O for (Ec.m.-VB)$\geq$9 MeV. The
$\langle$$\sigma_{F}$/$\sigma_{I}$$\rangle$ in this energy interval is
1.11$\pm$0.04 for 18O and 1.06$\pm$0.07 for 19F. This comparison between 18O
and 19F is complicated by the presence of the unpaired proton in 19F.

Examination of the reduced cross-section for 20Ne is more telling. For all
energies measured, the 20Ne reduced cross-section is lower than that of both
18O and 19F. Unfortunately, lack of data between $\sim$5
MeV$\leq$(Ec.m.-VB)$<$15 MeV prevents a better characterization of this
excitation function. The reduced cross-sections for 17F observed in this work
are reduced as compared to the cross-sections for 19F for the entire energy
range measured. This decreased value for 17F as compared to 19F is
qualitatively consistent with the trend observed for 18O and 16O.

To investigate the interaction of the valence protons on the valence neutrons
we have performed Relativistic Mean Field (RMF) calculations of the neutron
and proton density distributions of light nuclei. In examining the proton and
neutron density distributions for neutron-rich carbon nuclei from these RMF
calculations, one observes that with increasing neutron number not only does
the tail of the neutron density distribution extend further out but despite a
constant number of protons, the proton density distribution is slightly
extended [29]. This interaction between valence protons and neutrons is also
reflected in the one proton separation energies of 16O,18O,and 20O which are
12.1, 15.9, and 19.3 MeV respectively.

Based upon these calculations we hypothesize that the slight decrease in the
cross-section for 19F as compared to 18O could reflect the attraction of the
sd proton on the two sd neutrons resulting in a reduction of the fusion
probability. The lower cross-sections for 20Ne as compared to 18O and 19F is
similarly consistent with the attraction of the two sd protons on the two sd
neutrons and the consequent suppression of fusion. It should be appreciated
that the explanation of these observations is at a qualitative level and a
more quantitative description will require more sophisticated theoretical
calculations which include nuclear structure.

These observations for the fusion of light sd nuclei can be summarized as
follows. The fusion cross-section is enhanced by the presence of sd neutrons.
This increased cross-section is suppressed by the presence of sd protons. It
should be stressed that these changes in the reduced fusion cross-section
relative to the barrier are beyond the systematic changes expected. It is
clear that further experimental data, particularly in the case of 20Ne and 17F
would be extremely useful. To assess the sensitivity of our results to the use
of the interaction cross-section, $\sigma_{I}$, as a reference, we have
calculated the reduced cross-section as
$\sigma_{F}$/(AP${}^{\frac{1}{3}}$+AT${}^{\frac{1}{3}}$)2 where AP and AT are
the mass number of the projectile and target nuclei respectively. While more
sophisticated prescriptions for the reduced cross-section have been utilized
in comparing dissimilar systems, given the similarity of the systems compared
in this work this simple scaling is appropriate [30]. Examining the dependence
of this quantity on (Ec.m.-VB) in Fig. 4c reveals that although the reduced
cross-section for 20Ne lies closer to that of 19F, the trends observed in Fig.
4b remain the same indicating the robustness of the observations.

Figure 5: Comparison of the fusion excitation function for 17,19F+12C with the
predictions of a TDHF model (panel a) and CCFULL (panel b).

To investigate the role of dynamics in the 17,19F + 12C reaction we performed
time-dependent Hartree Fock calculations. On general grounds the TDHF approach
is well-suited to describing the large-amplitude collective motion associated
with fusion. Artificial symmetry restrictions are eliminated by performing the
TDHF calculations on a 3D cartesian grid [31]. Recent calculations [32] have
provided a good description of above-barrier fusion data [33, 34]. The TDHF
calculations were performed by initiating collisions for increasing impact
parameters in steps of 0.01 fm until the maximum impact parameter for fusion
is reached. Calculations were performed using the Sky3D model [35], with a SV-
bas interaction with DDDI pairing. It should be noted that pairing
significantly impacts the fusion cross-section in TDHF [15] emphasizing the
need for its accurate description. The results of the TDHF calculations for
17F and 19F are depicted in Fig. 5a as the solid and dashed lines
respectively. While the model provides a reasonable description of the 19F
data, the 17F cross-sections are overpredicted. It should be noted that TDHF
calculations often slightly overpredict the above-barrier cross-sections due
to their neglect of breakup processes, complicating direct comparison of the
cross-section. Nonetheless, it is noteworthy that the difference between 19F
and 17F observed experimentally is not reproduced.

Coupled-channels calculations [36, 37] have been performed to assess the
relative importance of the ground and excited states on the fusion cross-
section. Results of CCFULL calculations [38] in which only the ground state of
both target and projectile nuclei is considered are presented in Fig. 5b. A
reasonable description of both excitation functions is achieved. Inclusion of
excited states acts to increase the cross-section slightly, though primarily
at sub-barrier energies. The limited impact of excited states on fusion of 17F
in this near-barrier regime is consistent with other measurements [25, 7].
This description of the measured cross-sections corresponds to a potential in
CCFULL with RB= 7.89 fm, VB=9.17 MeV, and $\hbar$$\omega$=3.36 MeV for 17F and
RB= 8.42 fm, VB=8.58 MeV, and $\hbar$$\omega$=2.94 MeV for 19F. While the
limited span of the experimental data does not permit a reliable extraction of
$\hbar$$\omega$, the trend of increasing RB and decreasing VB with the
addition of two neutrons is qualitatively understandable reflecting the
increased nuclear interaction.

##  4 Summary

Systematic comparison of the fusion excitation functions for isotopes of O, F,
and Ne nuclei with a carbon target can be used to examine the impact of
valence sd protons and neutrons on fusion. After accounting for differences in
the static size of the incident nuclei and systematic changes in the fusion
barrier, differences are noted between fusion of 18O, 19F, 17F, and 20Ne with
a carbon target. At energies about 5 MeV above the barrier the reduced cross-
section for 17F decreases as compared to 19F. This decrease is similar to the
one observed for 16O as compared to 18O. The observed trend can be interpreted
as the influence of sd protons and neutrons on the reduced cross-section. In
this framework, at energies upto 20 MeV above the barrier the presence of
valence sd neutrons acts to increase the effective size of the system that
fuses above the increase in the static size. Both the 16O/18O and 17F/19F
manifest this behavior. Comparison of a nucleus with sd protons and neutrons
to one with sd neutrons alone indicates that the presence of the sd protons
results in a decrease of the effective size. This behavior is interpreted as
the strong interaction between the sd protons and the sd neutrons. It should
be emphasized that this explanation of the trends noted is qualitative and a
quantitative description which requires an accurate description of pairing
effects is not yet available. Both a widely-used analytic model of fusion and
a state-of-the-art dynamical model were used to investigate the systematic
behavior expected. While both models predict a decrease in the cross-section
with removal of sd neutrons, the magnitude of the observed reduction in the
fusion cross-section for 17F as compared to 19F is not reproduced. Coupled-
channels calculations which consider only the ground state nuclei are able to
reproduce the measured excitation functions despite the presence of the low-
lying 2s1/2 proton-halo state. This initial observation of the sensitive
interplay of valence neutrons and protons in the fusion of sd shell nuclei
motivates further high-quality measurements of fusion for neutron-rich light
nuclei. A new generation of radioactive beam facilities [3, 2], and in
particular the availability of low-energy reaccelerated beams, provides an
unprecedented opportunity to explore this topic and improve our understanding
of low-density nuclear matter.

##  5 Acknowledgements

This work was supported by the U.S. Department of Energy under Grant Nos. DE-
FG02-88ER-40404 (Indiana University), and the National Science Foundation
under Grant No PHY-1491574 (Florida State University). J.V. acknowledges the
support of a NSF Graduate Research Fellowship under Grant No. 1342962.

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11institutetext: Department of Physics and Astronomy, University of Turku,
Vesilinnantie 5, FI-20500, Finland
11email<EMAIL_ADDRESS>22institutetext: School of Physics, University College
Dublin, Belfield, Dublin 4, Ireland 33institutetext: INAF – Osservatorio
Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy
44institutetext: Yunnan Observatories, Chinese Academy of Sciences, Kunming
650216, P.R. China 55institutetext: Key Laboratory for the Structure and
Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650216,
P.R. China 66institutetext: International Centre of Supernovae, Yunnan Key
Laboratory, Kunming 650216, P.R. China 77institutetext: Dipartimento di Fisica
e Astronomia ”G. Galilei”, Università degli studi di Padova Vicolo
dell’Osservatorio 3, I-35122 Padova, Italy 88institutetext: School of
Sciences, European University Cyprus, Diogenes street, Engomi, 1516 Nicosia,
Cyprus 99institutetext: European Southern Observatory, Alonso de Córdova 3107,
Casilla 19, Santiago, Chile 1010institutetext: Institut d’Estudis Espacials de
Catalunya (IEEC), Gran Capità, 2-4, Edifici Nexus, Desp. 201, E-08034
Barcelona, Spain 1111institutetext: Institute of Space Sciences (ICE, CSIC),
Campus UAB, Carrer de Can Magrans, s/n, E-08193 Barcelona, Spain
1212institutetext: INAF - Osservatorio Astronomico di Brera, Via E. Bianchi
46, I-23807, Merate (LC), Italy 1313institutetext: Oskar Klein Centre,
Department of Astronomy, Stockholm University, Albanova University Centre,
SE-106 91 Stockholm, Sweden 1414institutetext: Department of Astronomy,
University of California, Berkeley, CA 94720-3411, USA 1515institutetext:
Physics Department, Tsinghua University, Beijing, 100084, China
1616institutetext: Beijing Planetarium, Beijing Academy of Science and
Technology, Beijing, 100044, China

# Red eminence: The intermediate-luminosity red transient AT 2022fnm

S. Moran  0000-0001-5221-0243 11 R. Kotak  0000-0001-5455-3653 11 M. Fraser
0000-0003-2191-1674 22 A. Pastorello 33 Y.-Z. Cai  0000-0002-7714-493X 44 5 5
6 6 G. Valerin 77 S. Mattila 1188 E. Cappellaro 33 T. Kravtsov
0000-0003-0955-9102 9911 C. P. Gutiérrez 10101111 N. Elias-Rosa 33 11 11 A.
Reguitti  0000-0003-4254-2724 1212 3 3 P. Lundqvist  0000-0002-3664-8082 1313
A. V. Filippenko  0000-0003-3460-0103 1414 T. G. Brink  0000-0001-5955-2502
1414 X.-F. Wang 1515 16 16

(Received ?; accepted ?)

We present results from a five-month-long observing campaign of the unusual
transient AT 2022fnm, which displays properties common to both luminous red
novae (LRNe) and intermediate-luminosity red transients (ILRTs). Although its
photometric evolution is broadly consistent with that of LRNe, no second peak
is apparent in its light curve, and its spectral properties are more
reminiscent of ILRTs. It has a fairly rapid rise time of 5.3$\pm$1.5 d,
reaching a peak absolute magnitude of $-12.7\pm$0.1 (in the ATLAS $o$ band).
We find some evidence for circumstellar interaction, and a near-infrared
excess becomes apparent at approximately +100 d after discovery. We attribute
this to a dust echo. Finally, from an analytical diffusion toy model, we
attempted to reproduce the pseudo-bolometric light curve and find that a mass
of $\sim$4 M⊙ is needed. Overall, the characteristics of AT 2022fnm are
consistent with a weak stellar eruption or an explosion reminiscent of low-
energy type IIP supernovae, which is compatible with expectations for ILRTs.

##  1 Introduction

With the advent of wide-field surveys over the last two decades, new classes
or subgroups of transients have been emerging, necessitating a revision of our
understanding of the final stages of stellar evolution. One such grouping
occupies a range of luminosities between those typical of supernovae (SNe) and
classical novae, and are usually referred to as intermediate-luminosity
optical transients (Berger et al. 2009) or gap transients (e.g. Kasliwal,
2012). Here, we focus on these and the distinction between luminous red novae
(LRNe) and intermediate-luminosity red transients (ILRTs).

Figure 1: $B$-band image taken at the NOT (ALFOSC) at +3 d showing the
location of AT 2022fnm
($\alpha=12^{\mathrm{h}}25^{\mathrm{m}}33^{\mathrm{s}}.713$,
$\delta=+45^{\circ}41^{\prime}12^{\prime\prime}.45$, J2000) in its host
galaxy, NGC 4389.

Luminous red novae have a number of distinctive features. Ranging in
luminosity from $-$4 to $-$15 mag, they display double-peaked light curves
(e.g. Blagorodnova et al., 2017; Pastorello & Fraser, 2019; Pastorello et al.,
2019b). The initial peak is thought to result from the cooling emission of the
common envelope after its ejection in a stellar-merger scenario (Metzger &
Pejcha, 2017; Blagorodnova et al., 2021; Matsumoto & Metzger, 2022). Around
the time of the first peak, the spectra are rather blue and display strong
Balmer and Fe II emission lines (Cai et al., 2019, 2022b). Later, there is a
second peak that is often attributed to matter ejection or hydrogen
recombination (Pastorello & Fraser, 2019; Matsumoto & Metzger, 2022). By the
time of the second peak, the Balmer lines are weaker and a forest of metal
lines (mostly in absorption) appears (Pastorello et al., 2019b). At late
times, the spectra resemble those of an M-type star, displaying molecular
absorption bands such as TiO and VO (Martini et al., 1999; Lynch et al., 2004;
Pastorello & Fraser, 2019; Pastorello et al., 2021). The strongest evidence
for the stellar-merger origin of LRNe comes from V1309 Sco, which showed a
decrease in period consistent with an inspiralling binary (Tylenda et al.,
2011).

On the other hand, ILRTs usually rise to peak over about two weeks (though
this is based on a limited sample of about ten objects), and their peak
absolute magnitudes tend to fall between $-$11.5 and $-$14.5 (e.g. Cai et al.,
2018, 2021; Stritzinger et al., 2020b). They display a post-peak plateau,
somewhat reminiscent of type IIP SNe. ILRTs show calcium emission features at
all stages of their evolution: in particular, the [Ca ii]
$\lambda\lambda$7291,7323 doublet, characteristic of ILRTs, and the
$\lambda\lambda\lambda$8498,8542,8662 near-infrared (NIR) Ca ii triplet. There
are several proposed mechanisms for the production of ILRTs, including
electron-capture SNe from super-asymptotic giant branch star progenitors
(Thompson et al., 2009; Pumo et al., 2009), massive stars in dusty cocoons
experiencing outbursts (Bond et al., 2009; Humphreys et al., 2011), and
accretion onto a compact object (Kashi et al., 2010).

Figure 2: Multi-band light curve of AT 2022fnm. The inverted open triangles
represent upper limits. The vertical line marks the epoch of the onset of the
outburst. ATLAS data beyond +40 d are also shown with open symbols as they are
noisier. Figure 3: $r$-band comparison of AT 2022fnm with other ILRTs and LRNe
from the literature. All objects have been corrected for reddening with the
values listed in Table 2 and respective references. The reference epoch for
each object has been set to its first peak.

Luminous red novae and ILRTs can be said to belong to the same broad family of
red transient phenomena, and as such it can be difficult to disentangle the
two observationally. For example, the [Ca ii] $\lambda\lambda$7291,7323
doublet, characteristic of ILRTs, is also seen in some LRNe (Cai et al.,
2019). Furthermore, the forest of metal lines characteristic of LRNe can also
be seen in ILRTs (e.g. NGC 300-2008-OT1; Humphreys et al. 2011). The ILRTs SN
2008S and NGC 300-2008-OT1 each had a detected progenitor that likely did not
survive the explosive event (Botticella et al., 2009; Adams et al., 2016;
Valerin et al., in prep.), supporting the idea that such events are terminal.

AT 2022fnm ($\alpha=12^{\mathrm{h}}25^{\mathrm{m}}33^{\mathrm{s}}.713$,
$\delta=+45^{\circ}41^{\prime}12^{\prime\prime}.45$, J2000; see Fig. 1) was
first discovered on 31 March 2022 (UTC dates are used throughout this paper)
by the Asteroid Terrestrial-impact Last Alert System (ATLAS) survey (Tonry et
al., 2018; Smith et al., 2020a) and was given the internal survey designation
ATLAS22kez (Tonry et al., 2022). It was classified on 2 April 2022 as a
luminous blue variable by Bostroem et al. (2022) as part of the Global
Supernova Project (Howell & Global Supernova Project, 2017; Howell, 2019).

The region of the sky containing AT 2022fnm was well sampled by the ATLAS
survey, which allowed us to constrain the epoch of the onset of the outburst
to within a 3 d window. We set the outburst epoch as MJD 59667.98 $\pm$ 1.5
based on the midpoint between the first ATLAS detection (MJD 59669.47) and the
last non-detection (MJD 59666.50, with a limiting magnitude of 20.17 in the
ATLAS $c$ band).

AT 2022fnm is associated with the galaxy NGC 4389, for which there is a spread
of distances reported in the literature. We adopted a luminosity distance of
10.3 Mpc, as determined from the measured redshift of 0.00239 for a Hubble
constant of 69.6 km s-1 Mpc-1 (Bennett et al., 2014). This results in a
distance modulus of $\mu=30.1\pm 0.6$ mag. This distance falls roughly between
the recent Tully-Fisher distance from the Cosmicflows4 catalogue
($\mu=29.855\pm 0.40$ mag; Tully et al. 2013) and the kinematic distance after
correcting for local velocity fields, $\mu=30.70\pm 0.15$ mag, from the NASA
Extragalactic Database111http://ned.ipac.caltech.edu/(NED). We take AT 2022fnm
to have a foreground extinction of $A_{\mathrm{V}}=0.040$ mag from Schlafly &
Finkbeiner (2011, from NED), which assumes an $R_{\mathrm{V}}$ of 3.1 (Schultz
& Wiemer, 1975); motivated by the blue spectra and unremarkable colour
evolution, we assume negligible host-galaxy extinction.

This paper is organised as follows: in Section 2 we discuss the acquisition
and reduction of the optical and infrared data. In Section 3 we present the
photometric evolution of the object and compare it to other gap transients; a
full list can be found in Table 2. In Section 4 we address the spectroscopic
evolution. Section 5 contains our overall discussion. The data logs, both
photometric and spectroscopic, along with the details of the objects used for
comparison can be found in Appendix A.

##  2 Observational data

###  2.1 Optical and infrared imaging

We obtained optical imaging from a number of telescope and instrument
configurations: Nordic Optical Telescope (NOT) + ALFOSC (as part of the NUTS2
programme222https://nuts.sn.ie/), Liverpool Telescope (LT) + IO:O, Gran
Telescopio Canarias (GTC) + OSIRIS, Asiago Schmidt 67/92 + Moravian, and the
Las Cumbres Observatory Global Telescope (LCOGT) network of 1m telescopes. A
detailed log is given in Table LABEL:table:opt_phot. The images were reduced
in a standard manner, comprising trimming, bias subtraction and flat-field
correction. In the case of the NOT+ALFOSC and Asiago data, we performed these
reductions with the foscgui
pipeline333https://sngroup.oapd.inaf.it/foscgui.html developed by E.
Cappellaro. The IO:O data were reduced using the IO:O
pipeline444https://telescope.livjm.ac.uk/TelInst/Pipelines/\\#ioo. The Las
Cumbres Observatory data were automatically reduced by the banzai
pipeline555https://github.com/LCOGT/banzai. Imaging was also obtained from the
ATLAS and Zwicky Transient Facility (ZTF; Bellm et al. 2019) surveys. These
data are reduced automatically using bespoke pipelines (Smith et al., 2020b;
Masci et al., 2019; Magnier et al., 2016). All data were taken in Sloan ugriz,
Johnson-Cousins BV, and ATLAS co filter systems; the ATLAS o band is roughly
equivalent to r+i and ATLAS c to g+r.

We also obtained five epochs of NIR imaging using NOT+NOTCam; the log is given
in Table 4. The images were reduced using the NOTCam QUICKLOOK reduction
package666http://www.not.iac.es/instruments/notcam/guide/observe.html. A
master flat was created using bright and dim sky flats and bad pixels were
masked. The wide-field camera of NOTCam suffers significant optical distortion
for which we accounted using a distortion model. Sky images were created from
the dithered on-source exposures, which allowed us to subtract the dark
current along with the sky background. For both optical and NIR imaging, we
performed point-spread-function-fitting photometry using the SuperNOva
PhotometrY (SNOoPY) package777SNOoPY is a package for SN photometry using
point-spread-function fitting and/or template subtraction developed by E.
Cappellaro. A package description can be found at
http://sngroup.oapd.inaf.it/ecsnoopy.html.

###  2.2 Optical spectroscopy

We obtained ten epochs of long-slit optical spectra spanning +2 d to +90 d
from a number of different telescopes; the observation log can be seen in
Table 5.

We reduced the NOT+ALFOSC spectra using the foscgui data reduction pipeline,
whilst the Keck+LRIS spectrum was reduced using the LPipe data-reduction
pipeline Perley, 2019. The Telescopio Nazionale Galileo (TNG)+DOLORES and
GTC+OSIRIS spectra were reduced manually using standard IRAF888IRAF is a
software suite for the reduction and analysis of astronomical data. A
description can be found at https://iraf-community.github.io/ tasks. The
reduction process was similar for each instrument, with the two-dimensional
(2D) spectra trimmed, bias-subtracted, and flat-fielded before being extracted
to one dimension. Arc-lamp spectra were used to calibrate the wavelengths of
the spectra and a master sensitivity function was created from a standard-star
spectrum that was taken with an identical instrumental setup, before a
telluric correction was applied.

##  3 Photometric evolution

###  3.1 Light curve and colour

Figure 4: Colour evolution of AT 2022fnm and comparison objects. We calculated
the colours with observations taken within 0.3 d of each other except in the
case of $V-K_{\mathrm{s}}$, for which we allowed a $\pm$3.5 d window, owing to
the modest cadence of the NIR observations. Figure 5: Rise time and decline
rates plotted against the (first) $r$-band peak for AT 2022fnm and comparison
objects. Rise times plotted with filled markers are calculated against
explosion or outburst epoch, whilst those with unfilled markers are calculated
against the detection epoch. Rise times are calculated to the peak observed
magnitude, which could be more sensitive to photometric uncertainties for
objects whose light curve flattens around the peak.

The multi-band light curve of AT 2022fnm is shown in Fig. 2. It exhibits a
relatively smooth evolution in all bands, with a plateau between approximately
+15 and +50 d in the redder bands.

As the $o$-band limit is the closest to the discovery epoch, we used it to
calculate the rise time by fitting it with a second-order polynomial; the
resulting rise time is 5.3 $\pm$ 1.5 d, which is relatively well constrained
compared to the limited sample of ILRTs with secure rise times, and we find a
peak apparent magnitude of $17.36$ mag in $o$, corresponding to an absolute
magnitude of $-$12.7 $\pm$ 0.1. Overall, the light curve of AT 2022fnm is also
reminiscent of LRNe, though it only displays a relatively short plateau rather
than a second peak.

Figure 6: Pseudo-bolometric properties. Top panel: Pseudo-bolometric light
curve of AT 2022fnm. Second panel: Temperature versus time of a black body
that has been fitted to the pseudo-bolometric light curve. Third panel: Log of
the radius versus time of a black body that has been fitted to the pseudo-
bolometric light curve. Bottom panel: Expansion velocity (radius over time)
versus time of a black body that has been fitted to the pseudo-bolometric
light curve. Note that the radius is simply being divided by the elapsed time
at each point to generate this crude measure. The data have been corrected for
redshift and foreground extinction. Points corresponding to the SED fits from
Fig. 8 have been overplotted for comparison in the form of red triangles. The
figure has been truncated at +80 d as the light curve is dominated by
numerical artefacts beyond this epoch. Figure 7: Attempt to reproduce the
pseudo-bolometric light curve with a simple analytical diffusion model. All
three components of the model are shown along with their total, which fits the
observed luminosity well. The model is somewhat degenerate, and the details of
this are discussed in the text, as are the parameters. The best-fitting
explosion epoch from our modelling is MJD 59669.38, which is consistent with
our adopted outburst epoch (MJD 59667.98 $\pm$ 1.5).

We compare the $r$-band light curve of AT 2022fnm with ILRTs and LRNe in Fig.
3 (references in 2). For the first 100 d, the light-curve evolution of SN
2008S does not differ markedly from that of AT 2022fnm, but the rate of
decline is slower and it does not display the same clear peak and plateau that
is visible in the AT 2022fnm light curve. AT 2018aes (Cai et al., 2021) has a
similar peak to AT 2022fnm, but its behaviour over the first 40 d after peak
does not resemble AT 2022fnm at all, instead showing much slower evolution,
but after +40 d their decline rates are quite similar. PTF 10fqs, albeit
fainter than AT 2022fnm, shows a similar evolution, though it has a clear
plateau that lasts until about +70 d.

In Fig. 4 we show the colour evolution of AT 2022fnm in $B-V$, $g-r$, and
$V-K_{\mathrm{s}}$; all display a clear redward progression with time
following a brief blueward excursion in the first ten days for $B-V$ and $g-r$
(we do not have $V-K_{\mathrm{s}}$ data covering this period). AT 2021afy, a
LRN, shows qualitatively similar behaviour to AT 2022fnm in $B-V$, with a
particularly close match in evolution from about +40 d. On the other hand, SN
2008S, the prototypical ILRT, also shows similar evolution in this colour. The
limited data for AT 2022fnm in $V-K_{\mathrm{s}}$ make comparison more
difficult, but in general, we see a redward movement in both ILRTs and LRNe,
though the rate varies from object to object. SN 2008S, despite its slightly
less steep redward movement, is qualitatively the best match to AT 2022fnm.

In Fig. 5 we show both the rise time of AT 2022fnm and our comparison objects
and their decline rates over the first 15 d and over the period between +50 d
and +100 d (where applicable). The purpose of this plot was to determine where
AT 2022fnm fell in the parameter space and whether or not it showed a
particular kinship with either LRNe or ILRTs. At least in the case of our
limited sample, it seems that LRNe occupy a larger portion of the parameter
space both in terms of rise times and decline rates, whilst the ILRTs are
clustered together. This is consistent with the findings of Cai et al. 2022a.
AT 2022fnm occupies a region consistent with both LRNe and ILRTs.

###  3.2 Pseudo-bolometric light curve and spectral energy distribution
evolution

In Fig. 6 we show the pseudo-bolometric light curve of AT 2022fnm along with
evolution of the corresponding black-body temperature and radius. The light
curve was computed based on our $uBgcVroizJHK_{\mathrm{s}}$ data using the
superbol package.999https://github.com/mnicholl/superbol (Nicholl, 2018).
Owing to a gap in the data coverage, we only show the first 80 d (cf. Fig. 2).
After this point the interpolation between epochs and extrapolation across
different filters become unreliable, and the pseudo-bolometric light curve is
dominated by numerical artefacts. We see that the black-body radius first
decreases (around +5 d), before increasing again at around +15 d. Hydrogen
recombination could increase the transparency, but then another energy source
would be required to stop and reverse this behaviour. As such, a source of
ionising photons is required to push the photosphere back out; this will be
addressed in Section 5. The photospheric velocities evident in the bottom
panel of Fig. 6 are 4,000 km s-1 and higher in the first week, whilst the
spectra do not suggest velocities this high. However, the spectral lines in
emission form farther out in an optically thin region of the circumstellar
material (CSM), whereas the velocities derived from the black-body fitting
reflect an optically thick continuum that lies at a smaller radius.

Table 1: Analytical diffusion model parameters. $M_{\mathrm{env}}$ | $R_{\mathrm{env}}$ | $v_{\mathrm{env}}$ | $E_{\mathrm{env}}$ | $M_{\mathrm{bulk}}$ | $R_{\mathrm{bulk}}$ | $v_{\mathrm{bulk}}$ | $E_{\mathrm{bulk}}$ | $E_{\mathrm{tot}}$ | $M_{\mathrm{Ni}}$ | $t_{\mathrm{diff,^{56}Ni}}$
---|---|---|---|---|---|---|---|---|---|---
M⊙ | cm | km s-1 | erg | M⊙ | cm | km s-1 | erg | erg | M⊙ | days
0.15 | 1014 | 9000 | 4x1046 | 4 | 2x1013 | 5000 | 4.2x1048 | 4.24x1048 | 1.6x10-4 | 2.3

We attempted to reproduce the pseudo-bolometric light curve with a simple
analytical diffusion model that accounts for the luminosity of the cooling
ejecta and the 56Ni decay process. Details will be presented in Valerin et al.
(in prep.), but the fundamental design is based on Chatzopoulos et al. (2012)
and Arnett (1980, 1982). The resulting model light curve for AT 2022fnm is
shown in Fig. 7. To mimic both the sharp first peak as well as the subsequent
slow decline, we divided the ejecta into two regions. The outer layers (which
can be identified with the stellar envelope) are faster and less dense, and
therefore release their internal energy over a short timescale. On the other
hand, the inner parts of the ejecta contain the bulk of the mass and are
slower, resulting in a longer photon diffusion time. Hydrogen recombination,
though not directly treated, is accounted for in a rough manner by means of
the so-called diffusion time in the bulk material. The light curve is well
reproduced by the model, with both the peak and the decline captured.

We note that the mass and the velocity of the ejecta are coupled, which means
that an equivalent model could be contrived with a higher mass and velocity
or, equally, with a lower mass and velocity, so we cannot use this model to
constrain the energy of the outburst reliably. Nevertheless, we were able to
reproduce the pseudo-bolometric light curve with a plausible set of
parameters. The model considers the integrated radiated energy of the bulk
material and the envelope separately, but combined they give a total
integrated radiated energy of $E_{\mathrm{tot}}=4.24\times 10^{48}$ erg. The
best-fitting explosion epoch from our modelling is MJD 59669.38, which agrees
within the uncertainties with our adopted outburst epoch (MJD 59667.98 $\pm$
1.5). Overall, the model is consistent with a weak eruption or explosion of a
star of relatively low mass. The full set of model parameters can be found in
Table 1. More sophisticated modelling may be used in the future to provide
more robust explosion parameters.

In Fig. 8 we show a sequence of spectral energy distribution (SED) fits for AT
2022fnm. It can be seen that at late times the SED is best fit by two black
bodies, owing to the development of a NIR excess. The NIR photometric data are
sparse, though it does seem that the parameters of the second black body stay
fairly constant between +104 d and +120 d.

Figure 8: Sequence of SED fits for AT 2022fnm. At late times a second black-
body component (red) is required, and the sum of the two components is shown
in green.

##  4 Spectral evolution

The spectral sequence for AT 2022fnm is shown in Fig. 9. It initially displays
a strong blue continuum that weakens with time, and in the later epochs the
continuum clearly becomes red. There is a large decrease in relative
luminosity bluewards of H$\beta$ between the +13 d and +23 d spectra. The
development of such a red continuum is common to both ILRTs and LRNe.

Balmer lines dominate the spectra throughout the evolution of AT 2022fnm.
Balmer lines are seen in both ILRTs and LRNe, though at late times the lines
typically decrease in prominence significantly in LRNe. The complete evolution
of H$\alpha$ and H$\beta$ is plotted in Fig. 10 with the spectra having been
binned to a common sampling.

It can be seen in Fig. 10 that the H$\alpha$ emission feature narrows with
time and that H$\beta$ shows similar behaviour. An absorption line is present
bluewards of H$\alpha$ at 6485 Å in the +33 d spectrum, which we associate
with H$\alpha$. A similar absorption is present bluewards of H$\beta$ at 4802
Å. Both have velocities of approximately $-3500$ km s-1. In the +33 d
spectrum, two absorption features can be seen at 6574 Å and 6592 Å, just
redwards of H$\alpha$, and if they are associated with it, then they have
velocities of approximately $+500$ km s-1 and $+1300$ km s-1, respectively.

There is a noticeable H$\alpha$ absorption trough in the +2 d spectrum at
approximately 6559 Å, corresponding to a velocity of around $-150$ km s-1 and
with a full width at half maximum of approximately 240 km s-1. This feature is
not present in the +3 d or later spectra. If this feature is real, then it
must result from a thin shell of material in the immediate vicinity of the
outburst that was quickly overrun.

To assess the evolution of the Balmer lines, we fitted a Lorentzian profile at
each epoch to adequately capture the electron-scattering wings. We find that
over the first 15 d after the onset of the outburst, the full width at half
maximum velocity stays at approximately 625–750 km s-1 but drops monotonically
to around 400 km s-1 at +90 d. In Fig. 11 we compare one early (+8 d) and one
late spectrum (+90 d) of AT 2022fnm to spectra of other objects. In the upper-
left inset of each panel, we show a zoom-in of the H$\alpha$ region for each
object. In the upper-left inset of panel (a), the H$\alpha$ profile of AT
2022fnm resembles that of AT 2017be and SN 2008S in both shape and velocity,
though the peak of the latter appears to be slightly redshifted. In the upper-
left inset of panel (b), H$\alpha$ appears to be slightly blueshifted at late
times in the case of AT 2022fnm and AT 2017be, but slightly redshifted in the
case of SN 2008S and perhaps AT 2018hso, though the signal-to-noise ratio of
the line is not particularly good for that object. At early times all of the
objects (AT 2022fnm, SN 2008S, AT 2017be, AT 2018aes, and AT 2018hso) are
dominated by the Balmer lines, as is typical in ILRTs and LRNe. The overall
spectral evolution of AT 2022fnm resembles that of the ILRTs AT 2017be and AT
2018aes. The LRN AT 2018hso resembles AT 2022fnm at early times, but at late
times it displays a red continuum and there is little resemblance between the
two objects in terms of spectral lines, firmly placing AT 2022fnm in ILRT
territory, as far as its spectra are concerned.

Moving on from the Balmer lines to consider oxygen, we note that there is the
narrow absorption at around 7770 Å, which is visible in the +33 d spectrum as
well as the preceding spectra. If this line is O I $\lambda$7774, then it is
unusual to see it in absorption; it does not appear to be visible in galactic
emission in the 2D spectra, so it is unclear whether this is a result of
subtraction. Typically, the O I $\lambda$7774 line is seen in emission in
airglow spectra. O I emission at 8446 Å is apparent by +57 d, growing to much
greater prominence by +90 d. This O I emission is not pronounced in SN 2008S
or AT 2017be at early times, but is clear at later times in these objects, as
can be seen in Fig. 11.

It is difficult to confirm the presence of Fe II above the noise in the +2 d
spectrum, but it is clearly present from the +3 d spectrum all the way up
through the +90 d spectrum. In particular, there are emission lines at 4924 Å,
5018 Å, and 5169 Å, as well as a forest of emission lines between
approximately 4490 Å and 4620 Å.

Turning to calcium, initially weak, we note that the Ca ii
$\lambda$8498,8542,8662 NIR triplet becomes apparent within the first two
weeks and is extremely pronounced at late times. The triplet is also weak at
early times in AT 2017be (Fig. 11), but, as in AT 2022fnm, it is very
prominent at later times. The [Ca ii] $\lambda\lambda$7291,7323 doublet,
characteristic of ILRTs, though present from the +13 d spectrum onwards, is
unusually weak for an ILRT. Like AT 2022fnm, the late-time spectrum of the
ILRT AT 2017be does not show particularly prominent emission in the [Ca ii]
$\lambda\lambda$7291,7323 doublet (Fig. 11). The doublet, a key signature of
ILRTs, is not typically seen in LRNe, although it was observed in AT 2018hso
as a very weak emission feature (Cai et al., 2019). The Ca H $\lambda$3969 and
Ca K $\lambda$3934 lines are initially very prominent in absorption, but by
+57 d they are instead in emission. The change from Ca H and K absorption to
emission can be seen in Fig. 11 where Ca H and K emission is also seen in the
ILRTs SN 2008S and AT 2018aes as well as in the LRN AT 2018hso.

The late-time spectra are reminiscent of interacting SNe, although we note
that many CSM-dominated transients appear similar at these phases. We see no
evidence for spectral features associated with nucleosynthesis (e.g. [O ii]
$\lambda\lambda$6300,6364), although these can be suppressed in interacting
core-collapse SNe due to higher densities. O i $\lambda$8446 is present, and
it is possible that this is due to Bowen fluorescence (Bowen, 1947). In this
case, the Lyman-$\beta$ photons required presumably arise from CSM
interaction. Given that the Ca ii $\lambda$8498,8542,8662 NIR triplet and that
lines due to Fe ii are present, whilst the [Ca ii] $\lambda\lambda$7291,7323
doublet is very weak (Fig. 9) and [Fe ii] lines are absent, suggests that the
density remained high (Kozma & Fransson, 1998; Dessart & Hillier, 2011).

Figure 9: Optical spectral sequence of AT 2022fnm. The spectra have been
calibrated against the photometry. The offsets applied to each spectrum are
shown in the plot, and the +90 d spectrum has been multiplied by a factor of 2
for clarity. Negative fluxes have been set to zero. The +2 d spectrum is
plotted in grey with a version binned to 10 $\AA$ overplotted in black to
compensate for the noise of the spectrum. The absorption features seen at
approximately 6725 Å and 8200 Å are artefacts. Figure 10: Evolution of
H$\alpha$ and H$\beta$ in AT 2022fnm over time. All spectra have been binned
to a common pixel size and normalised against the peak of H$\alpha$ in the
left panel and H$\beta$ in the right. Only spectra with both H$\alpha$ and
H$\beta$ coverage are shown. Vertical lines marking $\pm 1000$, 2000, and 3000
km s-1 have been added to guide the eye. Figure 11: Comparison of an early
spectrum and late spectrum of AT 2022fnm with other objects. The spectra have
been normalised against the sum of the flux densities except in the case of
the insets, which have been normalised against the peak of H$\alpha$. The
phases are calculated relative to the epoch of the onset of the outburst. The
upper-left inset shows H$\alpha$. The +83 d AT 2018hso spectrum has been
trimmed to 4500–9000 Å and binned to 10 $\AA$, owing to noise.

##  5 Discussion

The similarity of the light curve of the ILRT AT 2022fnm with light curves of
LRNe reflects that in both cases a comparable amount of energy is injected
into a comparable amount of hydrogen-rich material. The object displays a
clear NIR excess (apparent by approximately +100 d), typical of ILRTs (Cai et
al., 2021) and also seen in a number of LRNe (Pastorello et al., 2019a;
Stritzinger et al., 2020a). On the other hand, the object does not display the
double-peaked light curve associated with LRNe, though this feature is not
particularly prominent in some LRNe (e.g. AT 2018hso; Cai et al. 2019). Nor
does AT 2022fnm display the overwhelming forest of metal lines in absorption
along with the weakening of the Balmer lines typically seen at intermediate
times when LRNe have their second peak, or the molecular lines associated with
LRNe at late times. In certain respects, AT 2022fnm resembles AT 2018hso,
though AT 2018hso is a LRN with a double-peaked light curve.

Unlike LRNe, which display a second peak in photospheric radius after the
maximum, the photospheric radii of ILRTs typically undergo a monotonic
decrease (Botticella et al., 2009; Cai et al., 2019, 2021). AT 2022fnm, on the
other hand, displays aberrant behaviour in this regard, which precludes the
use of this metric as a tidy method of separating LRNe and ILRTs. In Fig. 6
over the first 20 d we see the photosphere moving inwards and then outwards
again. The decrease in radius could potentially be explained by increasing
transparency due to hydrogen recombination, which is indeed seen at around
6000 K in SN IIP models, for example (Kasen & Woosley, 2009).

After the decrease in radius, a source of ionising photons would then be
required to push the photosphere outwards. The spectral evolution provides
some clues: the continuum towards the blue end weakens between +13 d and +23
d, which is potentially indicative of increased opacity as the bluer
wavelengths experience greater scattering. Furthermore, the beginning of the
increase in radius at around +15 d also corresponds to the beginning of the
plateau seen in the light curve. Taken together, these observations could
naïvely be seen to suggest increased interaction, with the outer layers of
ejecta hitting the CSM and leading to the creation of photons and hence
ionisation, increasing the opacity. Alternatively, we may be seeing a scenario
similar to the increase in opacity routinely observed in SNe IIP, where the
shock ionises the outer envelope, leading to a decrease in transparency. The
radius of AT 2022fnm shows an overall variation of approximately a factor of
1.6, whilst the low-luminosity classical IIP SNe SN 2005cs and SN 2009md
display variation of approximately a factor of 2.5 (Jäger et al., 2020).
However, in contrast to AT 2022fnm, the early evolution of SN 2005cs shows a
monotonic increase in radius over time (unfortunately, no early radius
measurements are available for SN 2009md; Jäger et al., 2020).

The increase in radius in AT 2022fnm from around +15 d is puzzling given that
no temperature rise corresponding to the proposed ionisation is seen, though
the decline rate of the luminosity does slow at this epoch. Additionally, the
equivalent width of H$\alpha$ (plotted in Fig. 12) is decreasing at this
point. Similar, though less marked, behaviour was seen in the case of the ILRT
SNhunt120 (Stritzinger et al., 2020b), whose photospheric radius also dipped
early before climbing again.

Figure 12: Evolution of the spectral flux density and the equivalent width of
H$\alpha$. Top panel: Evolution of the spectral flux density of H$\alpha$ over
time. Middle panel: Evolution of the equivalent width of H$\alpha$ over time.
Bottom panel: Evolution of the ratio of H$\alpha$ to H$\beta$ spectral flux
density over time.

It is unclear whether the requirement of a second black-body component after
around +100 d is due to the presence of multiple ejecta components or the
presence of pre-existing dust. If there is pre-existing dust, it must not
exist in the immediate vicinity of AT 2023fnm or it would be quickly overrun
by the ejecta. The second black-body component has a temperature of $2200\pm
350$ K, consistent with the sublimation temperature for amorphous carbon dust
of approximately 1900 K (Landt et al., 2023). Unfortunately, the point at
which a second black-body component becomes necessary is approximately
coincident with the epoch of our final spectrum, and the lack of later spectra
makes its difficult to make any further determination.

Spectroscopically, AT 2022fnm more closely resembles an ILRT than a LRN.
Whilst the development of a red continuum is common to both LRNe and ILRTs,
the presence of the [Ca ii] $\lambda\lambda$7291,7323 doublet is a key
signature of ILRTs, though, as noted, it has also been seen in the LRN AT
2018hso. Ca H and K emission is visible in the +57 d and +90 d spectra, and in
earlier spectra it is visible in absorption. However, these lines are also
seen in emission in some LRNe and ILRTs (e.g. AT 2018hso and SN 2008S). The Ca
ii $\lambda$8498,8542,8662 NIR triplet is also prominent in both ILRTs and
LRNe, as seen in Fig. 11. Nevertheless, AT 2022fnm does not develop the forest
of metal lines typically seen in absorption in LRNe, nor does it develop the
strong molecular lines associated with LRNe at late times. This highlights the
need for late-time spectra to disentangle these two classes of objects.

The absorption lines bluewards of H$\alpha$ and H$\beta$ (at 6485 Å and 4802
Å, respectively), which we identify as components of these lines, have a
velocity of about $-3500$ km s-1. This is likely associated with the ejecta as
the photosphere recedes in mass coordinates, revealing further material.
Indeed, this recession in mass coordinates is apparent in Fig. 6 from around
+30 d, when the photospheric radius stays roughly constant. There are also
further absorption features redwards of H$\alpha$ at 6574 Å and 6592 Å, with
velocities of 500 km s-1 and 1300 km s-1, respectively.

AT 2022fnm can be said to be an ambiguous object in that it shows features
consistent with LRNe and ILRTs, straddling the phase space covered by both
categories of objects. AT 2022fnm is not alone in this respect, with certain
other objects, such as AT 2018hso, sharing features of both the LRN and ILRT
classifications (Cai et al., 2019). The simple analytical modelling we
performed shows that the object is consistent with a weak eruption or
explosion of a star of relatively low mass, though, as mentioned, the
parameters of the model are highly degenerate. The presence of pre-existing
dust, as evidenced by the NIR excess, is consistent with an ILRT; however,
both pre-existing and newly formed dust can also be found in core-collapse
SNe, such as SN 1987A (Roche et al., 1989; Lucy et al., 1989) and the type IIP
SNe SN 2002hh (Meikle et al., 2006) and SN 2004et (Kotak et al., 2009), as
well as in stellar-merger scenarios (Bond et al., 2003; Pastorello et al.,
2019a), so it cannot be used to discriminate between these possibilities.
Nevertheless, the preponderance of evidence supports the classification of AT
2022fnm as an ILRT, and the characteristics of the object are consistent with
a weak stellar explosion or eruption. We expect the next generation of surveys
to yield hundreds of LRNe and ILRTs, and suggest that a combination of
relatively high-cadence optical and NIR photometry and spectroscopy will be
needed if we wish to attribute various processes associated with the evolution
of single and massive binary systems.

###### Acknowledgements.

We thank the anonymous referee for carefully reading the manuscript. S. Moran
acknowledges support from the Magnus Ehrnrooth Foundation and the Vilho, Yrjö
and Kalle Väisälä Foundation. RK acknowledges support from the Research
Council of Finland (340613). MF is supported by a Royal Society – Science
Foundation Ireland University Research Fellowship. AP, EC and NER and AR
acknowledge support of the PRIN-INAF 2022 project “Shedding light on the
nature of gap transients: from the observations to the models”. Y.-Z. Cai is
supported by the National Natural Science Foundation of China (NSFC, Grant
12303054), the Yunnan Fundamental Research Projects (Grant 202401AU070063) and
the International Centre of Supernovae, Yunnan Key Laboratory (Grant
202302AN360001). S. Mattila acknowledges support from the Academy of Finland
project 350458. CPG acknowledges financial support from the Secretary of
Universities and Research (Government of Catalonia) and by the Horizon 2020
Research and Innovation Programme of the European Union under the Marie
Skłodowska-Curie and the Beatriu de Pinós 2021 BP 00168 programme, from the
Spanish Ministerio de Ciencia e Innovación (MCIN) and the Agencia Estatal de
Investigación (AEI) 10.13039/501100011033 under the PID2020-115253GA-I00
HOSTFLOWS project, and the programme Unidad de Excelencia María de Maeztu
CEX2020-001058-M. NER acknowledges partial support from MIUR, PRIN 2017 (grant
20179ZF5KS) and from the Spanish MICINN grant PID2019-108709GB-I00 and FEDER
funds. PL wishes to thank the Swedish Research Council for support. AVF’s
group at UC Berkeley has received financial assistance from the Christopher R.
Redlich Fund, Gary and Cynthia Bengier, Clark and Sharon Winslow, Alan Eustace
(W.Z. is a Bengier-Winslow-Eustace Specialist in Astronomy), William Draper,
Timothy and Melissa Draper, Briggs and Kathleen Wood, Sanford Robertson
(T.G.B. is a Draper-Wood-Robertson Specialist in Astronomy), and numerous
other donors; additional support was provide by NASA/HST grant AR-14295 from
the Space Telescope Science Institute (STScI), which is operated by the
Association of Universities for Research in Astronomy (AURA), Inc., under NASA
contract NAS5-26555. This work is supported by the National Natural Science
Foundation of China (NSFC grants 12033003 and 11633002), the Scholar Programme
of Beijing Academy of Science and Technology (DZ:BS202002) and the Tencent
Xplorer Prize. Based in part on observations made with the Nordic Optical
Telescope, owned in collaboration by the University of Turku and Aarhus
University, and operated jointly by Aarhus University, the University of Turku
and the University of Oslo, representing Denmark, Finland and Norway, the
University of Iceland and Stockholm University at the Observatorio del Roque
de los Muchachos, La Palma, Spain, of the Instituto de Astrofisica de
Canarias. The data presented here were obtained in part with ALFOSC, which is
provided by the Instituto de Astrofisica de Andalucia (IAA) under a joint
agreement with the University of Copenhagen and NOT. Based in part on
observations made with the Gran Telescopio Canarias (GTC), installed in the
Spanish Observatorio del Roque de los Muchachos of the Instituto de
Astrofísica de Canarias, in the island of La Palma. The W. M. Keck Observatory
is operated as a scientific partnership among the California Institute of
Technology, the University of California and NASA; the observatory was made
possible by the generous financial support of the W. M. Keck Foundation. We
thank WeiKang Zheng and Yi Yang for assistance with the Keck observations.
This work has made use of data from the Asteroid Terrestrial-impact Last Alert
System (ATLAS) project. ATLAS is primarily funded to search for near-Earth
objects through NASA grants NN12AR55G, 80NSSC18K0284, and 80NSSC18K1575;
byproducts of the NEO search include images and catalogues from the survey
area. The ATLAS science products have been made possible through the
contributions of the University of Hawaii Institute for Astronomy, the Queen’s
University Belfast, the Space Telescope Science Institute, and the South
African Astronomical Observatory. This work makes use of observations from the
LCOGT network. We acknowledge ESA Gaia, DPAC, and the Photometric Science
Alerts Team.

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##  Appendix A Tables

Table 2: Comparison objects. Name | $z$ | Type | Peak ($r$ or $R$) |  $A_{\mathrm{r}}$ or $A_{\mathrm{R}}$ | Rise | Decline (0$-$15 d, 50$-$100 d) |  Source(s)
---|---|---|---|---|---|---|---
|  |  |  |  | days | mag/day ($r$ or $R$) |
SN 2008S | 0.0002 | ILRT | -13.26 | 0.740 (R) | 7.0 | 0.0185, 0.0378 |  Thompson et al. 2009; Adams et al. 2016
NGC 300-2008-OT1 | 0.00048 | ILRT | -13.37 | 0.027 (R) | 28.3 | 0.0118, 0.0298 |  Thompson et al. 2009; Humphreys et al. 2011; Adams et al. 2016
PTF 10fqs | 0.008 | ILRT | -11.929 | 0.089 (r) | 10.80 | 0.0535, 0.0332 |  Kasliwal et al. 2011
AT 2013la | 0.002712 | ILRT | -13.24 | 0.024 (r) | 7.0 | 0.0353, 0.0162 |  Cai et al. 2021
AT 2017be | 0.001438 | ILRT | -12.66 | 0.123 (r) | 10.5 | 0.0329, 0.0182 |  Cai et al. 2018
AT 2017jfs | 0.008 | LRN | -14.12 | 0.055 (r) | 2.5 | 0.0689, -0.0068 |  Pastorello et al. 2019a
AT 2018aes | 0.00391 | ILRT | -12.70 | 0.053 (r) | 26.7 | 0.0181, 0.0660 |  Cai et al. 2021
AT 2018bwo | 0.00156 | LRN | -11.18 | 0.048 (r) | 13.2 | 0.0388, - |  Blagorodnova et al. 2021; Pastorello et al. 2023
AT 2018hso | 0.0039 | LRN | -13.06 | 0.026 (r) | 10.0 | 0.0740, -0.0005 |  Cai et al. 2019
AT 2019zhd | -0.001764 | LRN | -9.67 | 0.142 (r) | 61.3 | 0.0818, - |  Pastorello et al. 2021
AT 2021afy | 0.007208 | LRN | -12.63 | 0.130 (r) | 6.0 | 0.0167, 0.0415 |  Pastorello et al. 2023
AT 2021biy | 0.002035 | LRN | -13.08 | 0.039 (r) | 11.53 | 0.0666, -0.0177 |  Cai et al. 2022b
Table 3: Continued.

UTC Date | MJD | Epoch (d) |  $u$ (err) |  $B$ (err) |  $V$ (err) |  $g$ (err) |  $r$ (err) |  $i$ (err) |  $z$ (err) | Telescope (Instrument)
---|---|---|---|---|---|---|---|---|---|---
2022-04-01 | 59670.4 | 2.4 | 19.08 (0.39) | - | - | - | - | - | - | LCOGT (fa05)
2022-04-01 | 59670.5 | 2.5 | - | 18.44 (0.27) | 17.96 (0.27) | - | - | - | - | LCOGT (fa05)
2022-04-01 | 59670.6 | 2.6 | - | 18.42 (0.14) | 17.93 (0.08) | 18.18 (0.05) | 17.74 (0.05) | 17.65 (0.16) | - | LCOGT (kb82)
2022-04-02 | 59671.0 | 3.0 | 18.7 (0.04) | - | - | - | - | - | - | LCOGT (fa11)
2022-04-02 | 59671.1 | 3.1 | - | - | - | - | 17.59 (0.04) | - | - | NOT (ALFOSC)
2022-04-02 | 59671.2 | 3.2 | 18.69 (0.11) | - | - | - | - | - | - | LCOGT (fa05)
2022-04-02 | 59671.4 | 3.4 | - | - | - | 17.9 (0.04) | - | - | - | ZTF
2022-04-02 | 59671.5 | 3.5 | - | - | - | 17.91 (0.02) | - | - | - | LCOGT (fa05)
2022-04-03 | 59672.1 | 4.1 | - | 18.11 (0.02) | 17.61 (0.03) | 17.88 (0.02) | 17.49 (0.03) | 17.44 (0.02) | - | LCOGT (fa05)
2022-04-05 | 59674.4 | 6.4 | - | - | - | 17.92 (0.06) | - | - | - | ZTF
2022-04-05 | 59674.5 | 6.5 | - | - | - | - | 17.61 (0.01) | - | - | ZTF
2022-04-06 | 59675.9 | 7.9 | - | - | - | - | 17.77 (0.04) | - | - | NOT (ALFOSC)
2022-04-07 | 59676.2 | 8.2 | 18.56 (0.05) | 18.27 (0.03) | 17.87 (0.02) | 18.03 (0.02) | 17.87 (0.03) | 17.74 (0.02) | 17.78 (0.04) | LT (IO:O)
2022-04-07 | 59676.3 | 8.3 | - | - | - | 18.02 (0.08) | - | - | - | ZTF
2022-04-07 | 59676.4 | 8.4 | - | - | - | - | 17.89 (0.11) | - | - | ZTF
2022-04-07 | 59676.9 | 8.9 | - | - | 17.98 (0.1) | 18.17 (0.18) | 17.95 (0.07) | 17.83 (0.1) | - | Asiago Schmidt (Moravian)
2022-04-08 | 59677.0 | 9.0 | 18.71 (0.14) | 18.35 (0.09) | - | - | - | - | - | Asiago Schmidt (Moravian)
2022-04-09 | 59678.2 | 10.2 | 18.88 (0.06) | 18.49 (0.03) | 18.08 (0.02) | 18.28 (0.02) | 17.99 (0.02) | 17.96 (0.02) | 18.0 (0.05) | LT (IO:O)
2022-04-09 | 59678.3 | 10.3 | - | - | - | 18.27 (0.05) | - | - | - | ZTF
2022-04-11 | 59680.2 | 12.2 | 19.21 (0.1) | 18.70 (0.06) | 18.25 (0.03) | 18.49 (0.05) | 18.15 (0.03) | 18.18 (0.03) | 18.16 (0.05) | LT (IO:O)
2022-04-12 | 59681.0 | 13.0 | - | - | - | - | 18.20 (0.09) | - | - | NOT (ALFOSC)
2022-04-13 | 59682.2 | 14.2 | - | - | - | 18.59 (0.14) | - | - | - | ZTF
2022-04-13 | 59682.3 | 14.3 | - | - | - | - | 18.29 (0.09) | - | - | ZTF
2022-04-15 | 59684.3 | 16.3 | - | - | - | 18.71 (0.1) | 18.32 (0.08) | - | - | ZTF
2022-04-18 | 59687.3 | 19.3 | - | - | - | - | 18.40 (0.04) | - | - | ZTF
2022-04-18 | 59687.8 | 19.8 | - | - | - | 18.88 (0.13) | - | - | - | Asiago Schmidt (Moravian)
2022-04-18 | 59687.9 | 19.9 | - | - | - | - | 18.40 (0.14) | 18.34 (0.13) | - | Asiago Schmidt (Moravian)
2022-04-20 | 59689.2 | 21.2 | - | 19.26 (0.05) | 18.58 (0.02) | 18.97 (0.04) | 18.45 (0.02) | 18.33 (0.02) | 18.33 (0.04) | NOT (ALFOSC)
2022-04-20 | 59689.3 | 21.3 | - | - | - | 19.0 (0.11) | 18.48 (0.06) | - | - | ZTF
2022-04-21 | 59690.9 | 22.9 | 20.19 (0.17) | 19.40 (0.06) | 18.68 (0.05) | 19.06 (0.08) | 18.52 (0.05) | 18.37 (0.03) | 18.35 (0.04) | LT (IO:O)
2022-04-23 | 59692.9 | 24.9 | 20.43 (0.15) | 19.56 (0.07) | 18.73 (0.05) | 19.09 (0.06) | 18.56 (0.06) | 18.40 (0.03) | 18.36 (0.04) | LT (IO:O)
2022-04-24 | 59693.2 | 25.2 | - | - | - | 19.11 (0.12) | - | - | - | ZTF
2022-04-24 | 59693.3 | 25.3 | - | - | - | - | 18.58 (0.03) | - | - | ZTF
2022-04-26 | 59695.3 | 27.3 | - | - | - | 19.13 (0.12) | - | - | - | ZTF
2022-04-27 | 59696.9 | 28.9 | - | 19.75 (0.12) | 18.79 (0.07) | 19.17 (0.04) | 18.59 (0.03) | 18.43 (0.03) | 18.32 (0.03) | LT (IO:O)
2022-04-28 | 59697.3 | 29.3 | - | - | - | 19.18 (0.06) | 18.60 (0.07) | - | - | ZTF
2022-04-29 | 59698.0 | 30.0 | - | - | - | - | 18.61 (0.17) | - | - | NOT (ALFOSC)
2022-04-30 | 59699.2 | 31.2 | - | - | - | 19.18 (0.05) | - | - | - | ZTF
2022-05-02 | 59701.1 | 33.1 | - | - | - | - | 18.67 (0.05) | - | - | NOT (ALFOSC)
2022-05-02 | 59701.3 | 33.3 | - | - | - | 19.18 (0.11) | - | - | - | ZTF
2022-05-02 | 59701.9 | 33.9 | 20.84 (0.18) | 19.87 (0.05) | 18.88 (0.04) | 19.31 (0.04) | 18.67 (0.04) | 18.44 (0.05) | 18.34 (0.03) | LT (IO:O)
2022-05-06 | 59705.9 | 37.9 | - | 19.99 (0.06) | 19.0 (0.05) | 19.45 (0.05) | 18.73 (0.04) | 18.48 (0.02) | 18.35 (0.03) | LT (IO:O)
2022-05-10 | 59709.2 | 41.2 | - | - | - | - | 18.71 (0.09) | - | - | ZTF
2022-05-10 | 59709.3 | 41.3 | - | - | - | - | 18.77 (0.12) | - | - | ZTF
|  |  |  |  |  |  |  |  |  |
UTC Date | MJD | Epoch (d) | u (err) | B (err) | V (err) | g (err) | r (err) | i (err) | z (err) | Telescope (Instrument)
2022-05-12 | 59711.3 | 43.3 | - | - | - | - | 18.74 (0.13) | - | - | ZTF
2022-05-12 | 59711.9 | 43.9 | - | 20.18 (0.2) | 19.20 (0.06) | 19.66 (0.08) | 18.80 (0.04) | 18.59 (0.03) | 18.37 (0.04) | LT (IO:O)
2022-05-16 | 59715.2 | 47.2 | - | - | - | 19.71 (0.13) | - | - | - | ZTF
2022-05-19 | 59718.2 | 50.2 | - | - | - | - | 18.97 (0.1) | - | - | ZTF
2022-05-20 | 59719.9 | 51.9 | - | 20.45 (0.09) | 19.38 (0.05) | 19.93 (0.05) | 19.04 (0.04) | 18.74 (0.05) | 18.54 (0.04) | LT (IO:O)
2022-05-24 | 59723.9 | 55.9 | - | 20.69 (0.17) | 19.61 (0.13) | 20.15 (0.09) | 19.25 (0.04) | 18.91 (0.05) | 18.74 (0.1) | LT (IO:O)
2022-05-25 | 59724.9 | 56.9 | - | - | - | - | 19.29 (0.04) | - | - | GTC (OSIRIS)
2022-06-02 | 59732.0 | 64.0 | - | 21.04 (0.09) | 19.95 (0.04) | 20.41 (0.05) | 19.53 (0.03) | 19.18 (0.03) | 18.92 (0.06) | LT (IO:O)
2022-06-09 | 59739.9 | 71.9 | - | 21.67 (0.37) | 20.44 (0.16) | 20.98 (0.12) | 19.93 (0.08) | 19.52 (0.04) | 19.15 (0.07) | LT (IO:O)
2022-06-17 | 59747.0 | 79.0 | - | 22.34 (0.18) | 20.83 (0.06) | 21.45 (0.14) | 20.30 (0.04) | 19.89 (0.03) | 19.48 (0.04) | NOT (ALFOSC)
2022-06-18 | 59748.0 | 80.0 | 23.42 (0.47) | - | - | - | - | - | - | NOT (ALFOSC)
2022-06-22 | 59752.0 | 84.0 | - | 22.68 (0.22) | 21.19 (0.11) | 21.79 (0.1) | 20.59 (0.09) | 20.05 (0.05) | 19.72 (0.07) | NOT (ALFOSC)
2022-06-26 | 59756.9 | 88.9 | - | - | 21.56 (0.15) | 22.24 (0.26) | 20.87 (0.06) | 20.36 (0.05) | 20.10 (0.07) | NOT (ALFOSC)
2022-07-11 | 59771.9 | 103.9 | - | - | - | - | 21.70 (0.39) | - | - | LT (IO:O)
2022-07-12 | 59772.0 | 104.0 | - | - | - | - | - | 21.20 (0.22) | 20.75 (0.23) | LT (IO:O)
2022-07-13 | 59773.9 | 105.9 | - | - | 22.55 (0.24) | 23.42 (0.4) | 21.78 (0.17) | - | - | NOT (ALFOSC)
2022-07-18 | 59778.9 | 110.9 | - | - | - | - | 22.26 (0.18) | 21.41 (0.18) | 20.81 (0.17) | NOT (ALFOSC)
2022-07-31 | 59791.9 | 123.9 | - | - | - | - | - | 21.84 (0.44) | 21.04 (0.15) | NOT (ALFOSC)
2022-08-06 | 59797.9 | 129.9 | - | - | - | - | - | 22.20 (0.26) | 21.26 (0.17) | NOT (ALFOSC)
2022-08-15 | 59806.9 | 138.9 | - | - | - | - | - | 22.84 (0.46) | - | NOT (ALFOSC)

Table 4: NIR photometry given in Vega magnitudes. UTC Date | MJD | Epoch |  $J$ (err) |  $H$ (err) |  $K_{\mathrm{s}}$ (err) | Telescope (Instrument)
---|---|---|---|---|---|---
2022-05-07 | 59706.0 | 38.0 | 17.46 (0.04) | 17.11 (0.06) | 16.87 (0.04) | NOT (NOTCam)
2022-06-01 | 59731.0 | 63.0 | 17.84 (0.04) | 17.35 (0.06) | 17.02 (0.06) | NOT (NOTCam)
2022-07-09 | 59769.9 | 101.9 | 18.94 (0.09) | 17.84 (0.06) | 17.42 (0.09) | NOT (NOTCam)
2022-07-26 | 59786.9 | 118.9 | 19.16 (0.10) | 18.08 (0.09) | 17.59 (0.09) | NOT (NOTCam)
2022-08-17 | 59808.9 | 140.9 | - | - | 17.63 (0.12) | NOT (NOTCam)
Table 5: Optical spectroscopy. Instrumental resolutions were determined based on measurement of the [O i] $\lambda$5577.34 skyline except in the case of those marked with an asterisk, where reference values were used. UTC Date | MJD | Epoch (d) | Exp. (s) | Slit (′′) |  $\Delta\lambda$ (Å) | Telescope (Instrument, Grism)
---|---|---|---|---|---|---
2022-04-01 | 59670.4 | 2.4 | 3000 | 2.1 | 7.83* | APO 3.5m (KOSMOS)
2022-04-02 | 59671.2 | 3.2 | 3000 | 1.3 | 17.49 | NOT (ALFOSC, #4)
2022-04-06 | 59675.9 | 7.9 | 3000 | 1.0 | 12.87 | NOT (ALFOSC, #4)
2022-04-11 | 59681.0 | 13.0 | 2467 | 1.0 | 12.84 | NOT (ALFOSC, #4)
2022-04-22 | 59691.1 | 23.1 | 3600 | 1.0 | 9.37* | TNG (DOLORES, LR-B + LR-R)
2022-05-02 | 59701.1 | 33.1 | 3600 | 1.3 | 17.17 | NOT (ALFOSC, #4)
2022-05-25 | 59724.9 | 56.9 | 3600 | 1.0 | 9.75* | GTC (OSIRIS, R1000R)
2022-05-27 | 59726.1 | 58.9 | 3600 | 1.5 | 13.79* | TNG (DOLORES, LR-R)
2022-06-28 | 59758.3 | 90.3 | 3600 | 1.0 | 5.43* | Keck (LRIS, 600/4000 + 400/8500)


*[OMP]: orthogonal matching pursuit
  *[MP]: message passing
  *[MIM]: mutual-information-maximization
  *[BS]: base station
  *[ RF]: radio frequency
  *[SNR]: signal-to-noise ratio
  *[NMSE]: normalized mean square error