11institutetext: Nicolaus Copernicus Astronomical Center, Polish Academy of
Sciences, Rabiańska 8, 87-100 Toruń, Poland,
11email<EMAIL_ADDRESS>22institutetext: Nordic Optical Telescope,
Rambla José Ana Fernández Pérez 7, 38711 Breña Baja, Spain 33institutetext:
Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, 8000
Aarhus C, Denmark 44institutetext: Max-Planck-Institut für Radioastronomie,
Auf dem Hügel 69, 53-121 Bonn, Germany 55institutetext: Astronomical
Observatory, University of Warsaw, Al. Ujazdowskie 4, 00-478 Warszawa, Poland

# A radical transition in the post-main-sequence system U Equulei

Tomek Kamiński 11 Mirek R. Schmidt 11 Anlaug Amanda Djupvik 2233 Karl M.
Menten 44 Alex Kraus 44 Krystian Iłkiewicz 55 Thomas Steinmetz 11 Muhammad
Zain Mobeen 11 Ryszard Szczerba 11

###### Abstract

Context. U Equ is an unusual maser-hosting infrared source discovered in the
1990s. It was tentatively classified as a post-AGB star with a unique optical
spectrum displaying rare emission and absorption features from molecular gas
at a temperature of about 500 K. In 2022, we serendipitously discovered that
its optical spectrum has drastically changed since the last observations in
the 1990s.

Aims. We aim to characterize the drastic change in the spectrum and analyze
the photometric behavior of the object since 1989.

Methods. Optical high-resolution spectra of U Equ from the Southern African
Large Telescope are supplemented by archival data and near-infrared photometry
from the Nordic Optical Telescope. New spectral line observations with the
Effelsberg 100 m radio telescope and ALMA are presented. Radiative transfer
modeling of multiple epoch spectral energy distributions is performed.

Results. No circumstellar molecular features are present in the contemporary
optical spectra of U Equ. Non-photospheric absorption and emission from
neutral and ionized species dominate the current spectrum. Some of the
observed features indicate an outflow with a projected terminal velocity of
215 km s-1. Broad H&K lines of [Ca II] indicate a photosphere of spectral type
F or similar. For the first time, we find SiO $J=$1–0 $\varv$=1 maser emission
in U Equ. Our collected photometric measurements show that the source has been
monotonically increasing its optical and near-IR fluxes since about the
beginning of this century and continues to do so. The current rise in the
optical regime is about 1 mag. Spectral energy distributions at different
epochs show the presence of dusty circumstellar material that is very likely
arranged in a highly-inclined disk. Adopting a distance of 4 kpc, informed by
the Gaia parallax of U Equ, we find that the source’s luminosity is on the
order of 104 L☉. This luminosity has likely increased by a factor of a few in
the last decades, which is most probably related to the drastic change in the
optical circumstellar spectrum of the object.

Conclusions. The object has changed considerably in the last three decades
either due to geometrical reconfiguration of the circumstellar medium,
evolutionary changes in the central star, or owing to an accretion event that
has started in the system very recently. Observationally, U Equ appears to
resemble the Category 0 of disk-hosting post-AGB stars of Kluska et al.
(2022), especially the post-common envelope binary HD 101584. It is uncertain
if the drastic spectral change and the associated optical/mid-IR rise in
brightness witnessed in U Equ are common in post-AGB stars but such a radical
change may be related to the real-time onset of the evolution of the system
into a planetary nebula. We find that the post-AGB star V576 Car has undergone
a similar transformation as U Equ in the last few decades, so the phenomenon
is not extremely rare.

###### Key Words.:

Stars: AGB and post-AGB – circumstellar matter – Stars: mass-loss – Stars:
peculiar – Stars: individual: U Equ

##  1 Introduction

The transition of an asymptotic giant branch (AGB) star into a white dwarf
(WD) is a relatively poorly known process. The post-AGB (or pre-planetary
nebula; PPN) phase is expected to last on the order of 10 to 105 years
(Schoenberner, 1983; Miller Bertolami, 2016), during which the star is thought
to shed most of its outer layers and expose them to the harsh radiation of its
inner core, which, with time, will cool down to become a proto-WD (Paczyński,
1971). Although this phase is short compared to the main-sequence and the AGB
timespan, the number of evolved AGB stars in the Galaxy is big enough (e.g.,
Iwanek et al., 2022) for us to be able to catch some objects in this
transition stage. The hallmark of this evolutionary shift, as predicted by the
theory, should be a steep increase in the effective temperature of the
photosphere, occurring at a nearly constant luminosity. However, so far, we
have not identified firmly such objects in our observations, especially at the
onset of the post-AGB evolution when the star is still embedded in a cloud
consisting of dust and atomic/molecular gas accumulated during the turbulent
and mass-reducing AGB phase. Although the post-AGB star with increasing
effective temperature ($T_{\rm eff}$) cannot yet fully photoionize this
circumstellar envelope (which only happens in the planetary nebula, PN, phase)
the evolution off the AGB should have a significant impact on that
environment. The effect on the circumstellar medium may be actually easier to
spot than the direct change in the photospheric temperature. In practice, the
circumstellar environments of stars classified as post-AGB display an enormous
level of complexity (cf. Sanchez Contreras et al., 2022), often caused by
interactions with a companion, so that genuine changes caused by the actual
evolution of the star may not be easily recognizable.

U Equ (Barnbaum et al., 1996) and V576 Car (Couch et al., 2003) must be
amongst the most bizarre objects classified as post-AGB stars. When observed
at optical wavelengths in the 1990s, their spectra displayed strong molecular
bands in absorption (in both objects) or in both emission and absorption (U
Equ). The bands arise in warm ($\lesssim$1000 K) molecular gas located
relatively close to the stars, which, along our line of sight, are strongly
obscured by dust. Despite the severe obscuration, the main stellar component
in both objects was classified as types F–K which are too hot to be typically
associated with circumstellar environments rich in molecules, especially close
to the stellar photospheres. Indeed, U Equ and V576 Car have long been
considered very unusual. We focus here especially on U Equ, which showed
molecular electronic bands of various metal oxides in emission, most
prominently from TiO and VO. This is a very rare spectroscopic feature, and
only a handful of other objects have been reported to possess such emission
(Lloyd Evans, 1997). Red nova remnants, such as V4332 Sgr, V1309 Sco, and V838
Mon, have shown this phenomenon (Kamiński et al., 2015a, b, 2009) for years to
decades after their outbursts. Several young stellar objects displayed such
molecular emission features, too, in relation to an accretion-driven outburst
(e.g., Hillenbrand et al., 2012) but it is typically a short-lived feature in
the optical spectra. The red supergiant VY CMa is the only known source whose
optical spectrum has molecular emission lines permanently in emission, at
least since its first spectroscopic observations (Wallerstein, 1971; Herbig,
1974; Phillips & Davis, 1987; Kamiński et al., 2013; Humphreys et al., 2019).
U Equ is sometimes singled out as an example of a non-eruptive star with
molecular emission bands.

Figure 1: Spectra of U Equ obtained in 1994 (top; extracted from Barnbaum et
al., 1996) and in 2022–2023 (bottom). The spectrum had completely changed the
appearance from being dominated by molecular absorption-emission bands and
requiring low excitation temperatures into a circumstellar spectrum of
partially ionized atomic gas at a much higher excitation temperature. In this
paper, we describe and try to interpret this dramatic change in U Equ. The
bottom spectrum was smoothed to match the resolution of the earlier
observations. The earth symbols mark positions of strongest telluric features.

To investigate the phenomenon of molecular line emission in different types of
stars, we recently included U Equ in a survey in which we acquire high-
resolution spectra with the HRS
instrument111https://astronomers.salt.ac.za/instruments/hrs/ at the Southern
African Large Telescope (SALT). To our great surprise, the optical spectra
acquired in 2022 and 2023 for U Equ showed a completely different appearance
from those described in the literature. This dramatic change is illustrated in
Fig. 1 where we compare the recent SALT spectrum to one obtained in 1994. All
molecular features have disappeared and were replaced by numerous atomic lines
in absorption and emission, and arising from gas of low ionization but at
temperatures much higher than required for the historical circumstellar
spectra of U Equ. Here, we investigate the nature and cause of this remarkable
transformation of U Equ. In Sect. 2, we review what has been known about U Equ
before our observations. In Sect. 3, we present photometric data that allow us
to understand the object’s behavior over time scales of decades. We analyze
and discuss these changes in Sect. 7. Constraints on the physical
characteristics of U Equ put by Gaia observations are briefly presented in
Sect. 4. Complementary observations of molecular species at millimeter and
radio wavelengths are presented in Sects. 5 and 6, respectively. The SALT
observations are described in Sect. 8 and the new spectra are presented and
analyzed in Sect. 9. We discuss the essence of the observed changes and nature
of U Equ in Sect. 10. Therein, in Sect. 10.4, we also mention observations of
the other post-AGB star rich in optical molecular line absorption, V576 Car,
whose spectrum and photometric history have undergone similar changes as U
Equ.

##  2 Prior observations and impressions on U Equ

Although it was first reported by Reinmuth (1925) in a list of variable stars,
the modern history of U Equ can be traced back to the late 1980s, when far-
infrared (FIR) source that had been detected by the Infrared Astronomical
Satellite (IRAS 20547+0247) was recognized as a host of OH and H2O masers. The
maser observations of U Equ are described separately in Sect. 2.1. In addition
to photometric measurements, IRAS took a spectrum of U Equ with the Low
Resolution Spectrograph in the 8–22 $\mu$m range (Group, 1997). The spectrum
was described as having a strong 10 $\mu$m silicate feature. As a bright
infrared (IR) source with IRAS fluxes (epoch 1983.5) of 45.6, 33.7, 9.98, and
2.7 Jy at 12, 25, 60, and 100 $\mu$m, respectively, it was often classified as
an OH/IR star. However, several authors noticed that its location well below
the Galactic plane, with $l=-$26$\aas@@fstack{\circ}$11, is rather unusual for
this class of stars and might indicate a halo object.

Barnbaum et al. (1996, hereafter BOM96) brought U Equ to spotlight when they
acquired first optical spectra of the object in 1994. Their medium-resolution
spectra showed a multitude of absorption and emission features, the majority
of which were assigned to electronic bands of TiO, VO, and AlO. While the
emission bands most likely were produced by fluorescence, the absorption
spectrum was non-photospheric, given the lines’ narrow width and low inferred
excitation temperatures, that is with $T_{\rm ex}$=700–1500 K (most of the
bands were from transitions related to the ground electronic level and low
vibrational states). Such a spectrum appeared very unusual to the observers at
that time (cf. Lloyd Evans, 1997).

Based on the observation of the 10 $\mu$m silicate feature in absorption and
on IRAS colors, BOM96 postulated that U Equ has a circumstellar disk seen
edge-on and that this dusty disk was obscuring the star along our line of
sight. The optical molecular features were interpreted as arising in the inner
parts of the circumstellar medium, but BOM96 did not locate them in the disk
itself.

Based on that data, little could be inferred on the photospheric spectrum of U
Equ, as it was strongly contaminated by the omnipresent circumstellar
features. Based on the highly contaminated spectral features present in the
1994 spectra, BOM96 proposed the central star to be a giant (luminosity class
III) with a spectral type of mid G to early K. This was mainly based on a few
absorption lines of hydrogen and of a few metals. We consider these estimates
highly uncertain, based on the nature and quality of their spectra. The radial
velocity of the photospheric features could not be measured reliably in the
spectra, too.

BOM96 also combined photometric observations from different epochs to
construct a spectral energy distribution (SED) of the system. While an FIR
continuum excess, traced by the emission in the IRAS bands, indicated warm
dust at a temperature of about 350 K, the intrinsic stellar spectrum and
extinction components are not easily readable from their data. With their
preferred spectral type, however, BOM96 estimate the circumstellar reddening,
$E(B-V)$, as 0.68 mag (which corresponds to a visual extinction, $A_{V}$ of
2.1 mag); it is larger than their adopted interstellar reddening of
$E(B-V)$=0.09 mag. Assuming a luminosity of 103 L☉, from the SED they obtained
a distance of 1.5 kpc. BOM96 concluded that a dust-obscured evolved giant was
the most likely identification of the star, especially given the presence of
OH and H2O masers. The high Galactic latitude of the object, in their opinion,
excluded a young stellar object.

Geballe et al. (2005) observed U Equ spectroscopically in the near-IR in
1997–2003. They found an absorption features of CO and H2O which arise at
temperatures of 500–1000 K, similar to the temperature of the gas probed by
the optical bands of metal oxides. Resolved lines had a full width half
maximum of 35 km s-1 and were centered at –91 km s-1 in the heliocentric rest
frame (–78 km s-1 in the LSR). A weak ro-vibrational absorption line
$\varv$=1–0 $S(1)$ of H2 was also observed at a similar velocity. Line ratios
between the 12CO to 13CO isotopologues suggest a 12C/13C ratio of 4, but due
to opacity effects it may be significantly higher. Indeed, the value of 4
would be very perplexing in an O-rich object. Geballe et al. (2005) estimated
column densities for the three molecular species. The values they obtain
(e.g., $2\times 10^{20}$ cm-2 for CO) appear to be too low compared to what
one would infer from the circumstellar extinction of $A_{V}=2.1$ mag assumed
by BOM96. One possible interpretation is an extremely low dust content in the
circumstellar material – which would be contradictory to the FIR excess – but
the uncertainties in both quantities are very large. The central velocity of
absorption lines was interpreted by Geballe et al. (2005) as representing that
of the stellar photosphere. However, we note that the width and the excitation
temperature characterizing these lines readily imply a circumstellar origin.

Several imaging experiments in the visual and infrared, including one with the
Hubble Space Telescope (HST) and Very Large Telescope (VLT), did not spatially
resolve the U Equ system (Barnbaum et al., 1996; Geballe et al., 2005;
Siódmiak et al., 2008; Lagadec et al., 2011).

Figure 2: Top left: Maser emission of OH and H2O extracted from the literature
data is compared to the CO $J$=3–2 profile obtained more recently with ALMA.
The spectra were acquired with different spectral resolutions and are shown
here normalized to the peak. Right: On a reduced LSR velocity axis, the
spectrum of the SiO $J$=1–0 $\varv$=1 maser (red) is compared to that of the
1667 MHz OH maser (observed in 1987, orange) and to the CO $J$=3–2 line
(purple). Bottom: Spectra of the water masers acquired recently with the
Effelsberg telescope. They are shown in the absolute intensity scale to
demonstrate the variability in line fluxes. The velocity in the heliocentric
rest frame is 13.1 km s-1 lower than in the LSR frame adopted here.

###  2.1 Masers in the OH/IR phase

Masers in U Equ has been long thought to be very unusual. The H2O and OH
masers have been observed in U Equ at least since 1985 (Zuckerman & Lo, 1987)
and have been known to be variable in intensity and velocity. While
variability is a common phenomenon for maser lines, that observed toward U Equ
was quite remarkable. The first observations of the H2O 22 GHz masers revealed
a single-peak spectral feature at $V_{\rm LSR}$=–75.0 km s-1. However, later
observations in 1995 showed two weaker features, at –79.9 and –90.0 km s-1,
demonstrating behavior untypical for a circumstellar maser source. The OH
masers at lower frequencies have been observed much more often. In May 1987,
the 1612 and 1667 MHz OH masers were observed at LSR velocities of –80.4 and
–75.7 km s-1 (Sivagnanam et al., 1990), but only the 1612 MHz feature was
clearly double. The OH and H2O lines showed little overlap in velocities, and
only one component had a similar velocity in both species. Two components of
1612 MHz OH masers were also detected in August 1985 by Chengalur et al.
(1993) at –79.4 and –75.3 km s-1 (LSR). Barnbaum et al. (1996) observed the OH
1667 MHz masers in 1995 and noticed profile changes on time scales of months.
No double structure, known from the earlier epochs, was found. Etoka & Le
Squeren (2004) observed both OH transitions around the same time (in 1994) and
found only single-peaked features, as well. In the 1994–1995 observations, the
OH profile was centered near –75.7 km s-1, with an error of about a few km s-1
(the errors are not indicated in the original papers). In those observations,
the OH masers were getting weaker. The 1612 MHz emission was not detected in
1995. In a survey by Lewis (2002), the same OH transition was targeted but not
detected in 1999–2001. This led B. M. Lewis to conclude that U Equ (= OH
051.3+26.1) is an example of a ‘dead OH/IR star’, a fate he predicted for all
OH/IR stars on timescales of $\sim$1700 yr. Based only on the OH maser
observations, it is clear that the object was experiencing a radical change at
least since the beginning of the 1990s. We compare the literature reports on
the spectral positions OH and H2O maser lines in Fig. 2.

No SiO masers had ever been reported in U Equ. Deguchi et al. (2010) put only
upper limits with an rms of 0.06 Jy for the $J$=1–0 rotational lines in the
$\varv$=1 and 2 vibrationally-excited states in 2001. Any earlier reports
would be valuable, but we were not able to find any. Post-AGB objects are
almost never associated with SiO masers (Engels, 2001), but note the prominent
case of the pre-planetary nebula OH231.8+4.2 containing a Mira component
(Morris et al., 1987).

##  3 Archival photometry

To understand the transition that occurred in the optical spectrum of U Equ,
we tried to reproduce its multiband light curves from all available sources.
Below, we summarize and critically evaluate the archival materials that we
found. The light curves are shown in Fig. 3

###  3.1 Visual

Heidelberg-Königstuhl plates reported by Reinmuth (1925) provide two
measurements of a source at the location of U Equ (the coordinates of his
source number 95 converted to J2000 are RA=20:57:16.31 and Dec=02:58:44.78,
that is within 2″ from the current position). Magnitudes of ¡15.5 in 1902, and
of 14.5 in 1904 and 1915 were measured. These measurements are from unfiltered
observations with blue photographic plates.

More Heidelberg-Königstuhl blue plates were scanned and assigned an
astrometric solution. From those available at the Heidelberg Digitized
Astronomical Plates service222https://dc.zah.uni-
heidelberg.de/lswscans/res/positions/q/form, we selected ones in which U Equ
and nearby fields stars are readily seen. Those digitalized plates are from
1902, 1904, 1915, and 1927, and do not show any considerable variation (within
photometric accuracy) in the brightness of U Equ compared to the field stars
(see Fig. 8). On those dates, and likely in the entire 1902–1927 period, the
star did not experience any sudden flux increase by several magnitudes. It is
reasonable to assume that the magnitudes stayed close to those reported by
Reinmuth, that is near 14.5 mag, in the first quarter of the 20th century.
Some variability was however noticed since the star received a variable star
designation.

Plates of the first Palomar Observatory-National Geographic Sky Survey (POSS)
provide two measurements in 1951, 16.05 mag on a red E plate and 21.77 mag on
a blue O plate. These magnitudes were measured by McMahon et al. (2000). The
second POSS survey yielded a red magnitude of 13.86 in Sept. 1990, a blue
magnitude of 15.88 in June 1993, and $I$=12.41 in August 1994 (Monet et al.,
2003). The same catalog specifies the first-epoch red magnitude as 15.28 and
not 16.05 mag. We interpret this discrepancy as a result of intrinsic
uncertainties in photometry from these plates. Based on these measurements,
the photographic magnitudes did not change by more than 1 mag between 1951 and
1990.

According to Geballe et al. (2005), the visual magnitude of U Equ in 1994 was
9. No source of this measurement is given, and we find it very unlikely the
star had risen to such a low magnitude, especially given the photometric
constraints from the second POSS survey. Geballe at al. give also a magnitude
13 for 1996, 1998, and 1999, but again with no details of how those
measurements were acquired. We include these photometric points in Fig. 3 but
consider them very doubtful.

A $V$ band light curve spanning from 2004 to 2016 was extracted from the
INTEGRAL Optical Monitoring Camera (IOMC) data333https://sdc.cab.inta-
csic.es/omc/secure/form_busqueda.jsp?resetForm=true. We utilized their $V3$
measurements.

We extracted another set of visual magnitudes from the Catalina
Survey444http://nesssi.cacr.caltech.edu/DataRelease/ (Drake et al., 2009) for
the period 2005–2013. These are unfiltered magnitudes very roughly transformed
to $V$.

Several measurements in the $grizy$ bands from 2009–2014 were found in the
detection catalog of Pan-STARRS1555https://catalogs.mast.stsci.edu/panstarrs/
(Chambers et al., 2016). In a similar time period, measurements were taken by
the AAVSO Photometric All-Sky Survey (APASS)666https://www.aavso.org/apass,
mainly in the $BVgri$ bands.

The ongoing Sky Patrol of the All-Sky Automated Survey for Supernovae
(ASASSN)777https://asas-sn.osu.edu/ has been observing U Equ in $Vg$ bands
since 2012 (Shappee et al., 2014; Kochanek et al., 2017). We include here data
obtained till June 2023.

Several single-epoch measurements in various bands were found. We list them
here for completeness, but not all are included in our analysis.

  * •

A Johnson $B$ magnitude of 15.75 mag was derived from the Ukrainian glass
archive (Andruk et al., 2016)888https://cdsarc.cds.unistra.fr/viz-
bin/cat/I/342#/description for July 1991. This value is close to that from the
blue plate POSS-II measurement of 1993.

  * •

Gaia DR3 measurements of $G$=12.730$\pm$0.004, $BP$=13.184$\pm$0.011, and
$RP$=12.091$\pm$0.010 mag. The exact epochs of these measurements are unknown.

  * •

The AAVSO archive has a single $V$ measurement of 12.5 mag for Nov. 2021.

  * •

Siódmiak et al. (2008) provided HST magnitudes of ¡14.94 and 14.12 in the
$F606W$ and $F814W$ filters based on observations from May 2003. The $F606W$
image was saturated.

  * •

Public data of the TESS mission are available from Aug. and Sept. 2022. We
extracted the data using the Lightkurve package (Lightkurve Collaboration et
al., 2018) and scaled the raw instrumental intensity to the corresponding
observational data from ASASSN. The TESS data have high cadence and show
better some short-term variability of the source.

Figure 3: Photometric evolution of U Equ.

Multiband visual colors are mainly available for the period 2010–2013. No
colors are essentially known prior to 2005. The archival plates, although they
only sparsely probe the photometric behavior, seem so indicate unchanged
magnitudes in the blue and red. If there were any variations, they must have
been lower than about 1 mag between 1904 and 1994. Of course, it is always
possible that significant changes occurred in the large windows when no one
observed U Equ.

Measurement uncertainties were rarely given in the above quoted sources. The
combined photometric data are shown in Fig. 3.

###  3.2 Near-infrared

A collection of $JHK$ measurements between 1999 and 2005 were taken from the
Optical and Near Infrared Atlas of the Arecibo Sample of OH/IR
Stars999https://arecibo.cab.inta-csic.es/ (Jiménez-Esteban et al., 2021).
Hereafter and in Fig. 3, we call these points the Arecibo data.

Fouque et al. (1992) performed $JHKLM$ photometry of the object in Aug. 1989
using the ESO 1 m telescope.

BOM96 presented measurements in the $J$ and $K_{s}$ bands for Oct. 1995 and
obtained with the Shane telescope.

The 2MASS measurements in $JHK_{S}$ were obtained in July 2000 (Cutri et al.,
2003).

Lagadec et al. (2011) used the VISIR instrument on the ESO Very Large
Telescope (VLT) on 29 June 2008 to image U Equ at 8.59, 11.85, 12.81 $\mu$m
that is in the PAH1, SiC, and NeII filters, respectively. We transformed the
measured flux densities of 21.3, 35.7, and 41.7 Jy, respectively, into Vega
system magnitudes of 0.94, –0.27, –0.63 in those respective filters.

Observations in all four WISE bands (i.e., at $W1$=3.4, $W2$=4.6, $W3$=12,
and, $W4$=22 $\mu$m) were extracted from the AllWISE Multiepoch Photometry
Table101010https://irsa.ipac.caltech.edu/cgi-bin/Gator/nph-
scan?submit=Select&projshort=WISE (Wright et al., 2010). These observations
were made in May and Nov. 2010.

A much longer time sequence of photometric measurements at 3.4 and 4.6 $\mu$m
was extracted from the NEOWISE Reactivation Single Exposure (L1b) Source Table
(Mainzer et al., 2011). The $W1$ and $W2$ data were taken over the period from
2014 May 2014 to Oct. 2022. For better clarity, in Fig. 3 we rebinned the data
by averaging measurements obtained within 7 days (original data are shown in
gray).

The most recent $JHK_{S}$ photometry is not archival, but was obtained by
ourselves on 25 April 2023 with the NOTCam instrument (Abbott et al., 2000) at
the Nordic Optical Telescope (NOT; Djupvik & Andersen, 2010). The seeing
during the observation was of about 0$\aas@@fstack{\prime\prime}$8\. Owing to
the relatively high brightness of the target, the high-resolution (HR) camera
of NOTCam (0$\aas@@fstack{\prime\prime}$079/pixel) was used and exposure times
kept short: 30 s for the $J$-band, 10.8 s for the $H$-band, and 3.6 s for the
$K_{S}$-band. Moreover, for the $K_{S}$-band, the telescope was slightly
defocused to avoid saturation. The images were obtained by dithering in a
3$\times$3 pattern, and reduced with the IRAF package notcam.cl to correct for
non-linearity, flat-field, background sky, and to finally align, shift, and
calculate median combined images. A nearby field with six 2MASS stars,
observed right after U Equ with the same set-up, was used for the photometric
calibration. We converted the 2MASS photometry of these stars to the Maunakea
Observatory (MKO) system using the transformations in Leggett et al. (2006).
Aperture photometry yielded $J$=10.37$\pm$0.01 mag, $H$=9.05$\pm$0.04 mag, and
$K$=7.00$\pm$0.03 mag, where the errors reflect the uncertainty in the
photometric calibration. The $JHK$ systems used in earlier epochs may not be
the same.

##  4 Constraints from Gaia mission

Gaia measured the U Equ’s trigonometric parallax of 0.1891$\pm$0.0539 mas and
a proper motion of ($\delta_{\rm RA},\delta_{\rm Dec}$)=(10.3512$\pm$0.0671,
$-10.5706\pm 0.0447$) mas yr-1 (all from DR3 Gaia Collaboration et al., 2022).
The parallax indicates a geometric distance of 4.8${}^{+1.1}_{-0.9}$ kpc and a
photogeometric distance of 3.3${}^{+1.2}_{-0.5}$ kpc (Bailer-Jones et al.,
2021a), more than 2–3 times larger than adopted by BOM96111111For definitions
of geometric and photogeometric distance determinations, see Bailer-Jones et
al. (2021b). We adopt here 4 kpc as the distance to U Equ, as a compromise
between the geometric and photogeometric method. Like for many other evolved
and potentially binary stars, formal errors in both methods may be
underestimated.

Then, the Galactic latitude of U Equ implies that it is located 1.9 kpc below
the plane of the Galaxy and its projected motion is mainly away from the
Galactic plane; the proper motion in the Galactic frame is
($\delta_{l},\delta_{b}$)=(–3.4, –14.4) mas yr-1. Adopting the radial velocity
of –78 km s-1 (Geballe et al., 2005), the Galactocentric
motion121212https://docs.astropy.org/en/stable/coordinates/galactocentric.html#coordinates-
galactocentric of ($\varv_{x},\varv_{y},\varv_{z}$)=(–65.0, 11.8, –282.5) km
s-1 shows a relatively high speed away from the Galactic center. This makes U
Equ very likely a near-halo object, and as such it may be also metal poor.
Other halo objects have been identified near the position of U Equ (e.g.
Malhan et al., 2018) and the Phlegethon stream is located nearby (Ibata et
al., 2021). The interstellar extinction toward U Equ is, however, very low.
Using the maps of Schlafly & Finkbeiner
(2011)131313https://irsa.ipac.caltech.edu/applications/DUST, we find the most
likely value $E(B-V)$=0.078 mag (or $A_{V}$=0.24 mag). BOM96 adopted a
similarly low reddening of 0.09 mag. A more recent 3D dust map of Green et al.
(2019) provides a slightly higher value, $E(B-V)$=0.115 mag (or $A_{V}$=0.36
mag).

In the 120 years between the Gaia observations and the first photographic
plates, U Equ should have moved by about 1$\aas@@fstack{\prime\prime}$5\. This
figure is consistent with the appearance of the plates, but their seeing is of
a similar order.

Gaia DR3 results for U Equ include a flux-calibrated spectrum in the range
3330–10200 Å at a low spectral resolution. In Sect. 9 and Fig. 5, we make use
of a combined BP/RP averaged spectrum.

##  5 ALMA observations

For reference, we also extracted observations of the CO 3–2 pure rotational
transition near 345.8 GHz. It was found in the archive of the Atacama Large
sub-Millimeter Array (ALMA). The data were obtained for the NESS collaboration
(PI P. Sciculna; Scicluna et al., 2022) on 10 Aug 2021 with the Atacama
Compact Array (ACA) at an angular resolution of
3$\aas@@fstack{\prime\prime}$5\. The data were processed with the CASA
pipeline using default calibration tables. Next, data for spectral range
covering the CO transition were imaged using tCLEAN with various visibility
weighting schemes. In Fig. 2, we present the source-averaged spectrum of CO
3–2. With a 302 s exposure, the spectrum has a modest signal-to-noise ratio
(S/N) and Fig. 2 shows a smoothed version.

The same line was observed by us with the Atacama Pathfinder Experiment 12 m
submillimeter telescope (APEX) with the FLASH instrument (Heyminck et al.,
2006) on 29 June 2012 and was detected at an even lower S/N of $\leq$10\. Its
intensity was close to that measured with ALMA (7.1$\pm$0.5 Jy km s-1), but
uncertainties are large. This however indicates that there were no major
changes in the CO feature since 2012.

The ALMA dataset shows a weak continuum source of 8.07$\pm$0.92 mJy (1$\sigma$
error). The continuum was integrated within 330.2–333.7 and 343.2–346.2 GHz
avoiding the CO 3–2 emission. The continuum flux corresponds to an average
wavelength of 885 $\mu$m.

##  6 Effelsberg observations

#### SiO maser

U Equ was observed with the 100 m Effelsberg radio telescope on 11 June 2023
using the S7mm Double Beam receiver covering 40–45 GHz. The data were acquired
and calibrated using standard procedures. After 47 min of observing time on
source, we detected a weak (0.24 K $T_{A}$) and narrow (FWHM=1.4 km s-1) line
of SiO $\varv$=1 $J$=1–0 at $V_{\rm LSR}$=$-75.0$ km s-1. Given the narrow
line width, it is readily a maser line. It is shown in Fig. 2. Although other
transitions (from SiO and other molecules) were covered in the same spectrum,
for example SiO $\varv$=0 $J$=1–0, they were not detected at a nominal rms
noise level of 25 mK ($T_{A}$).

#### H2O masers

The 100 m telescope also observed the water transition
$J_{Ka,Kc}$=$6_{1,6}$–$5_{2,3}$ at 22.2 GHz (or 1.35 cm) using the S14mm
Double Beam receiver and the WFFTS backend. On 2 Nov. 2022, an integration of
19.5 min resulted in an rms noise of 28 mK (in $T_{A}$) and displayed at least
3 emission lines centered at LSR velocities of $-88.2$, $-68.3$, and $-66.4$
km s-1; three other features are seen near $-93.6$, $-62.8$, and $-53.4$ km
s-1 at an S/N of 6–8. The stronger lines have FWHMs between 0.6–1.1 km s-1,
indicative of masers. The strongest component has a peak intensity of 8.7 K.
The observations were repeated on 14 Aug. 2023 with a total integration time
of 29.4 min and a resulting noise rms of 79 mK. The spectrum of U Equ showed
much stronger lines at this later epoch. The spectra are compared in Fig. 2.
In Aug 2023, four strong components are seen at $-89.6$, $-87.8$, $-68.7$, and
$-66.6$ km s-1. There is also a weaker line centered at –61.4 km s-1. The
strongest component reaches 36.4 K in the peak. The line widths are in the
range 0.6–0.8 km s-1 but the weakest feature appears broader with a FWHM of
1.9 km s-1. The strong variability in intensity and position of the H2maser
components is obvious between 2022 and 2023, but is even more striking when
compared to the literature data (Sec. 2.1). Some of the early observations
might not even cover the most redshifted components. Overall, the H2O maser
components appear over the velocity range spanning from about $-90.5$ to
$-60.0$ km s-1, with the midpoint near –75 km s-1, which is also the position
of the SiO maser line.

##  7 Photometric analysis

The collected archival data in Fig. 3 show that U Equ is currently the
brightest at visual and IR wavelengths since observations of this object
started in 1904. Although it is not straightforward to compare the multiepoch
data obtained in different photometric systems and without the knowledge of
proper measurement uncertainties, it is apparent, especially in the visual
bands, that the brightness of U Equ has been monotonically increasing since
about 2009. The $JHK$ measurements probe a long period (1989–2006 and 2023),
but unfortunately do not cover well the period of the gradual visual rise. The
NOT $JHK$ measurements from 2023 are 1.1–1.5 mag lower than in 1989,
indicative of an increased near-IR flux by a factor of 2.8–4.0. The steady
rise is also evident in the WISE and NEOWISE observations, but it is visibly
steeper in the $W1$ band, that is at shorter wavelengths.

The steady rise is best documented by combined observations obtained with
different visual filters, albeit it is not covered entirely by consecutive
observations in a single band. The visual rise is shown in more detail in Fig.
4. There, linear fits to the Catalina $V$ data after 2006 and to all ASASSN
$g$-band data demonstrate that the rise had been steeper at its onset and has
flattened at the time of the most recent observations. The total magnitude
change in the visual is probably $\gtrsim$1 mag. Although at a low photometric
accuracy the rise appears to be relatively steady, the TESS data show small-
amplitude flux variations of $\sim$5% on a timescale of days.

Figure 4: Rise in brightness of U Equ in visual bands. The blue inset shows
variability documented by arbitrarily scaled TESS data. The inset in the lower
right compares the transmission curves of the $V$ and $g$ filters.

###  7.1 SED evolution

In an attempt to investigate the nature of the photometric changes, we
selected three “effective” epochs with mid-IR data. Based on the data
presented in Sec. 3, we compiled the SEDs as follows.

#### Epoch 2023

The 2023 compilation is defined by the epoch of the NOT $JHK$ measurements on
25 Apr. 2023. We extrapolated the NEOWISE data to the same date by fitting
linear functions to NEOWISE data only. We assign those points large
uncertainties of 0.5 mag, which reflect the fit and the large scatter in the
input data. We selected ASASSN $g$-band measurements from two dates closest to
the date of NOT observations (22 and 29 May) and averaged them to obtain a
representative optical-wavelength value. Finally, we also add to the SED the
ALMA continuum measurement from 2021. Although it was obtained nearly two
years earlier, the submm flux is not expected to have changed much. Even with
a large uncertainty, its addition to the 2023 SED puts useful constraints on
the models attempting to explain the source energetics in the most recent
epoch.

#### Epoch 2010

For the defining moment of the 2010 SED, we select the epoch of WISE
observations in all four bands. Nearly simultaneous with the WISE measurements
(within a week), observations were made in the $yz$ bands as part of PAN-
STARRS and in $V$ within the Catalina survey. We also add to this epoch $JHK$
fluxes obtained from linear interpolations between the latest measurements
from the Arecibo survey and those from NOT. Again, we assign large
uncertainties to these data points.

#### Epoch 1989

The earliest SED worth considering is defined by the simultaneous $JHKLM$
measurements obtained with the ESO 1 m telescope in 1989. Although the IRAS
flux measurements were obtained a few years prior, we add them to the 1989
epoch too, assuming no significant changes in the mid- and far-IR.had occurred
over the relevant time period. We also include the magnitude from 1990
measured on the red POSS plate. The three SEDs are compared in Fig. 5.

Figure 5: Comparison of SEDs at three epochs and to the Gaia average DR3
spectrum. For clarity, errorbars are omitted and data points are connected by
lines. Connectors do not represent the actual shape of the SEDs. The ALMA
measurement is omitted for display reasons. Data were corrected for
interstellar extinction.

All data used in SEDs were corrected for the interstellar extinction, assuming
$A_{V}$=0.24 mag (see Sec. 4). In order to understand the nature of the
changes taking place in U Equ, we modeled these three SEDs using different
dust radiative transfer tools. First, we attempted to reproduce the SEDs with
models in which the central star is surrounded by a spherical shell of
silicate dust. We used DUSTY (Ivezic et al., 1999) model grids compiled by the
Dusty Evolved Star Kit (DESK) (Goldman, 2020). Best fits were searched with
the $\chi^{2}$ test. Although solutions were found, they were hardly
satisfactory. The spherical models with different dust opacities always
predict too strong fluxes at near-IR wavelengths, producing a spectral slope
that is not observed in U Equ. From this, we conclude that spherical models
are not adequate for reproducing the geometry of the U Equ system. Despite
these caveats, for a distance of 4 kpc, the models converge at luminosities of
8 700 (epoch 2010) to 55 000 (epoch 2023). Also, the estimated gas mass-loss
rates increased from 10-5 to 10-4 M☉/yr between 2010 and 2023. Differences in
the models likely reflect insufficient spectral coverage and inadequacy of the
models, rather than a physical change in the system.

We next used a much wider grid of circumstellar SEDs compiled in Robitaille
(2017a). Although the database was created mainly for young stellar objects
(Robitaille, 2017b), it is applicable to a wider range of circumstellar
environments, including evolved systems. A wide range of models exists in the
grid with different sets of structural components, including a star, a flared
passive disk with or without an inner hole, an envelope with or without
cavities, and an ambient medium (see Robitaille, 2017b, for a full overview).
We performed analysis only for models with 4–9 free parameters. The models
were compared to data points corrected for interstellar extinction. Without
any other constraints, the sedfitter search was typically indicating best
models with a giant star of effective temperature, $T_{\rm\star}$, from 10 000
to 30 000 K (B spectral types) and an extra extinction of $A_{V}$=2 to 3 mag.
These synthetic SEDs reproduce the observations very well. However, since both
$T_{\rm\star}$ and $A_{V}$ of these models disagree with spectroscopic results
(Sect. 9) and with our constraints on the ISM extinction (Sects. 4 and 9.2),
we discarded these models.

We then narrowed the search for best Robatille SEDs by limiting the effective
temperature to lower than 9000 K and allowing for an extra interstellar
reddening of $A_{V}\leq$0.1 mag (cf. Sect. 4). The best solutions are shown in
Fig. 9. The fits for a wide range of circumstellar configurations are not
reproducing the observations as well as the unconstrained sedfitter results,
but are still better than the spherical DUSTY models. For example, Table 1
presents parameters for models where the star is surrounded by a passive disk
extending out from the sublimation radius (grid sp--s-i) or by a transition-
type disk with an inner hole (grid sp--h-i). Although U Equ has been described
in the past as a system possessing a disk, ours are the first results of
actual SED fits encompassing a disk. Because none of the considered models
satisfactorily reproduce the observations, we critically examine some of the
model parameters.

Table 1: Best star-disk models reproducing the SEDs.

| No hole models | Disk with inner hole
---|---|---
| 1989 | 2010 | 2023 | 1989 | 2010 | 2023
extra $A_{V}$ (mag) | 0.100 | 0.030 | 0.022 | 0.100 | 0.100 | 0.097
Stellar radius (R☉) | 81.3 | 36.8 | 70.1 | 96.7 | 32.4 | 94.1
Stellar temperature (K) | 7200 | 8100 | 5700 | 6700 | 8200 | 6800
Luminosity at 4 kpc (103 L☉) | 33 | 7.5 | 19 | 43 | 5.9 | 41
Disk dust mass (10-3 M☉) | 2.7 | 12 | 1.0 | 0.48 | 32 | 1.4
Disk inner radius (AU) | $R_{\rm sub}$ | $R_{\rm sub}$ | $R_{\rm sub}$ | 3.0 | 37.2 | 1.3
Disk outer radius (AU) | 2922 | 173 | 599 | 468 | 601 | 299
Disk flaring power, $\beta$ | 1.002 | 1.070 | 1.118 | 1.119 | 1.085 | 1.033
Disk surface density power | –1.73 | –0.37 | –1.27 | –1.09 | –0.45 | –0.74
Disk scaleheight (AU) | 9.9 | 18.8 | 18.3 | 19.1 | 18.2 | 12.6
Disk inclination (°) | 71.2 | 63.5 | 59.1 | 61.6 | 64.3 | 71.3
141414Disk in models with no inner hole starts at a respective dust
sublimation radius, Rsub. The luminosity was calculated from the radius and
temperature using the Stefan-Boltzmann law without the Robatille’s scaling
factor (cf. Robitaille, 2017b).

  * •

Although some types of models give the best result for the highest allowed
reddening of $A_{V}$=0.1 mag, most of the best fits indicate a much lower
extra extinction with an average of 0.05 mag. Our assumed extinction of
$A_{V}$=0.24 mag is likely not far from the true value.

  * •

Across the different types of models, the best stellar temperature and radius
are in the ranges of 8200$\pm$860 K and 50$\pm$18 R☉. From the Stefan-
Boltzmann law, these indicate a luminosity of about 9800 L☉. The uncertainties
are too large to investigate how these stellar parameters were changing over
time.

  * •

The models are not particularly constraining on the properties of the
circumstellar matter. If a disk were the dominant component, Table 1 implies
that: its mass is on the order of 10-4 to 10-2 M☉; it extends from a few AU to
a few hundred AU; and has a high inclination, that is, 59°–71° (so it is seen
nearly edge-on).

The system can be characterized better only with more densely spaced
contemporaneous observations, spanning from blue optical to submillimeter
wavelengths.

##  8 Spectroscopy with SALT

We used the High Resolution Spectrograph (HRS; Crause et al., 2014) at the
Southern African Large Telescope (SALT) to observe U Equ on: 5, 12, 13 May
2002; 5, 8, and 9 Aug. 2022; 28 May and 30 June 2023. The typical exposure
time was 1862 s. The HRS is a dual-beam échelle spectrograph operating in two
arms that cover 370–550 and 550–890 nm. We used the high resolution mode which
afforded a resolution of about 65 000. Because HRS is a fiber-fed instrument
with a separate fiber observing the sky at a certain angular distance from the
science target, subtraction of sky emission lines is not always perfect, which
affected our spectra. The telluric lines are, however, easily recognizable
when spectra from different epochs are compared at the heliocentric rest
frame. Data were reduced and wavelength-calibrated using a HRS
pipeline151515https://astronomers.salt.ac.za/software/hrs-pipeline/ with
default settings (Kniazev et al., 2017). Échelle orders and spectra from
different arms were merged. Owing to the telescope design and mode of
operations, SALT observations are not calibrated in flux. The pipeline-
processed spectra were thus normalized to the continuum, that is, were divided
by high-order polynomials. Spectra presented here were shifted to the
heliocentric velocity frame. Normalized spectra from different dates were
averaged. Due to a significant difference in the heliocentric velocity
correction during different observing seasons, telluric features are severely
broadened in the combined spectra. No reddening correction was applied.

##  9 Analysis of the visual spectra

The full spectrum of U Equ averaged over all 2022–2023 exposures is shown in
Figs. 10–13. As mentioned, the spectrum has changed dramatically since its
observations in the 1990s. Instead of a plethora of molecular bands in
emission and in absorption, we observe mainly circumstellar features of atomic
gas of a low ionization. The atomic spectrum is very rich. We identified more
than 600 features, most of which are absorption lines, but pure emission lines
and composite P Cyg profiles are present as well. We identified the lines
using several catalogs, chiefly from the
NIST161616https://physics.nist.gov/PhysRefData/ASD/lines_form.html, which
provides an interface for grouping lines in multiplets. A list of identified
lines is available from the authors upon request, but the identification is
also graphically presented in Appendix C. Among the identified species are H
I, Fe i, Fe ii, Ti i, Ti ii, Cr I, Cr II, Sr ii, V I, Ca i, Ca ii, Ni I, Mg I,
Si I, Si II, O I, C I, N I, Ba II, Co I, K I, and Na I. We found no helium
lines. Below, we highlight main groups of lines visible in the spectrum.

  * •

The Balmer and Paschen series of H I are clearly recognizable in the spectra
of U Equ. The lowest Balmer transitions, including H$\alpha$ and H$\beta$,
have strong emission centered near a heliocentric velocity of –94 km s-1. Pure
absorption is seen in lines from higher energy levels in both series. Based on
the profiles, the absorption components are in most cases non-photospheric and
form in an outflow crossing the line of sight. Additionally, H$\alpha$ has a
complex emission-dominated profile with extended wings ($\approx$50 Å wide at
the base) which are very likely formed due to photon scattering on electrons
in the recombining plasma. Although H$\alpha$ is by far the strongest emission
feature in the entire spectrum, it contributes only 9% to the total flux in
the $R$ band. The H$\beta$ line has a classic P Cyg profile, but photospheric
absorption may be recognizable in the blue-shifted part.

  * •

The only other lines showing broad emission wings comparable to those of
H$\alpha$ are lines of the Ca II IR triplet (rest wavelengths at 8498, 8542,
and 8662 Å). These lines have complex profiles, with multiple overlapping
absorption components. Additionally, each line of the triplet is contaminated
by overlapping emission of the Paschen series.

  * •

The H&K lines of Ca II near 3950 Å are dominated by broad absorption
components. These are photospheric features somewhat contaminated in the line
cores by circumstellar emission. By comparing H I and Ca II H&K features to a
grid of stellar spectra with known spectral types, we find that a type close
to F6 best represents U Equ’s photosphere. However, all F types and early G
types of different luminosity classes match these noisy photospheric features
comparably well. In Fig. 6, spectra of U Equ and two other F-type stars are
compared. The stellar photospheric temperature is likely around 5700 K, but
the small number of photospheric features and a modest S/N make this number
very uncertain.

  * •

Spectral features of CNO elements are of special interest. Neutral carbon is
seen through the forbidden emission near 8725 Å. Other lines of C I may be
present but lack an unambiguous identification. No features of N I were found.
Neutral oxygen is traced with two [O I] emission lines near 5577 (weak), 6300,
and 6363 Å, and with three permitted multiplets in absorption. The triplets at
around 7775 and 8446 Å have their upper energy levels near 11 eV, while the
weaker multiplet near 6155 Å is from relatively highly excited (13 eV) states.
These features are dominated by absorption and, since lines of a given
multiplet blend, they form very distinct wide features, known to
spectroscopists mainly from spectra of post-AGB stars. Future studies may
attempt to derive relative abundances of the CNO elements in the circumstellar
gas using these CNO tracers.

  * •

By the number of individual transitions, the spectrum is dominated by lines of
Fe i–ii ($\approx$270 lines) and of Ti i–ii ($\approx$170 lines). (We find no
forbidden lines of Fe ii.) The presence of such rich spectra makes unique
identification of individual lines rather difficult. The Fe and Ti spectra
show primarily absorption at shorter wavelengths and pure emission in the red
part of our spectra, but mixed, P Cyg, profiles are also observed. Chromium,
Cr i–ii, has a similar spectrum. An analysis of the Fe I pure emission lines
under the assumption of local thermodynamic equilibrium (LTE) presented in
Appendix D suggests that the neutral gas is excited at a temperature of about
4200 K.

  * •

Several lines from the ground level are observed. Among them, the Na i
doublet, studied in more detail below, is the most complex and displays
several absorption components overlapping with weak emission. The K i optical
doublet is present as well, mainly in weak absorption, but is severely
contaminated by telluric features. Two resonance lines of Ba II at 4554 and
4934 Å are identified in absorption. Other clear resonance lines include
absorption lines of Ti I, Sr II, and Fe I. The resonance lines of Fe I have
double profiles.

Figure 6: Comparison of spectra of the F-type stars HD167278 (blue) and
G188$-2$22 (red) from XSL to the current spectrum of U Equ (black). U Equ is
likely to have a spectral type F as well.

Based on the described characteristics of the optical spectra, the
circumstellar medium of U Equ has a low ionization degree, but is certainly
more ionized than it appeared 28 years ago. Such a low-ionization
circumstellar spectrum is not uncommon among post-AGB stars. In Appendix C, we
compare the 2022–2023 average spectrum of U Equ to that of a post-AGB star HD
101584 (Olofsson et al., 2019). The spectrum of HD 101584 which we use was
extracted from the Xshooter Spectral
Library171717http://xsl.u-strasbg.fr/index.html (XSL, Chen et al., 2014). The
visual spectrum of HD 101584 has been well described in the literature, for
instance in Sivarani et al. (1999, 2001) and in Kipper (2005).181818Whether HD
101584 is a post-RGB or post-AGB star is a matter of debate and depends on the
object’s true luminosity/distance, but the post-AGB status of this F0
supergiant has been recently reaffirmed by Parthasarathy et al. (2022). The
spectra of U Equ and HD 101584 are remarkably similar in terms of the presence
of atomic features and their general profiles, that is, when a feature has a P
Cyg profile in U Equ, it often has this type of profile in HD 101584, too.
Some differences exist, partially related to the better quality of the
Xshooter spectrum. A few pure absorption and pure emission features can be
found in HD 101584, but not in U Equ. Some of the differences can be
attributed to the difference in the relative abundances of the CNO elements,
but a dedicated study would need to confirm this. There are other spectra of
well known post-AGB stars that are very similar to what is currently observed
in U Equ, including the spectrum of the protoplanetary nebula M1-92
(Minkowski’s Footprint) described in Arrieta et al. (2005), which has an F2
central star.

###  9.1 Kinematic components

As mentioned, H$\alpha$, H$\beta$, and the Ca II IR triplet contain an
absorption component centered at a heliocentric velocity –94$\pm$3 km s-1 and
with a FWHM of about 26$\pm$5 km s-1. It can be seen in Fig. 14. This
component extends below the local continuum level in several of these lines,
so it definitely represents an absorption component rather than a gap between
two emission components. At the same velocity appears the red edge of some
absorption lines, for instance, in the Paschen lines (Fig. 16) and in the Na I
doublet (Fig. 17) — but not in absorption lines of Si II that have the blue
edge at a few km s-1 higher velocity. Near –94 km s-1 is also the cross-over
point from emission to absorption for several P-Cygni profiles, for instance,
in the profile of K I (Fig. 17) (but not in the P-Cyg profiles of Fe II); the
same velocity is also exactly in between the two absorption components of some
of the Fe I lines, which display a double absorption feature (e.g., Fig. 17).
The heliocentric velocity of –94 km s-1 corresponds to an LSR velocity –81 km
s-1, which is close to the velocity of the blue component of the 1612 MHz OH
maser observed in 1987 and to the red component of the H2O maser observed in
1995 (see Sect. 2.1). Based on the inventory of optical lines where the
component is seen, it may represent the coolest circumstellar gas on the line
of sight toward the star. Hereafter, we call it component A.

Pure emission lines, like those of [C I] $\lambda$8727 and [O I]
$\lambda$6363, are centered at a heliocentric velocity of –87$\pm$2 km s-1, so
are redshifted with respect to the absorption component described above by
$\approx$7 km s-1. The center of the emission lines is also the crossover
point between absorption and emission in the P-Cyg lines of Fe II and Ti II
(Fig. 18); it also sets the red edge of absorption in Si II (Fig. 16).
Converted to an LSR velocity of $-74$ km s-1, it matches very well the
position of the red component of the OH masers and the position of the water
masers observed in 1986 (see Sect. 2.1); it matches also to the velocity of
the SiO maser discovered in U Equ in 2023 and the midpoint of the H2O maser
components detected with the Effelsberg telescope (Sect. 6). Hereafter, we
call it component B.

The CO line observed by ALMA is centered exactly between components A and B.
It is uncertain which (if any) of these represents the systemic velocity of U
Equ, but since mm-wave emission of CO is undoubtedly thermally excited and is
not subject to circumstellar extinction, it may best represent the “average”
systemic velocity. Unfortunately, the absorption lines that we interpret as
photospheric are too contaminated and of a too low S/N to provide the stellar
velocity. However, a cross-correlation of the average spectrum near the Ca II
H and K lines against a synthetic spectrum of a giant with an effective
temperature of 6000 K yields a heliocentric velocity of –96.8$\pm$24.8 km s-1.
The large uncertainty makes this result of limited use, but the stellar
velocity seems to be closer to component A. This is puzzling, as circumstellar
SiO masers are usually considered to be very good indicators of the stellar
velocity. Of course, there may be two or more dynamical centers of the system,
for instance two stars, which could account for such a complex kinematic
structure.

Many of the lines of neutral and ionized species contain a broad absorption
feature, which in some cases is combined with emission, forming classical
P-Cyg profiles. This is undoubtedly a signature of an outflow. In lines of
hydrogen of the Paschen series, Na I, Fe II, Si II, the two O I triplets – the
absorption component spans from approximately –85 down to nearly –300 km s-1
(see Figs. 14–18). The outflow has thus a maximum (projected) expansion
velocity of 215 km s-1. This value is comparable to the escape velocity of a
giant of about one solar mass and a radius of 5–10 R☉. Not all absorption
tracers reach such a high velocity. For instance, the main absorption
component in the resonance line of Fe I $\lambda$5110 extends only to an
expansion velocity of 35 km s-1 (Fig. 17). It would thus appear that the high-
velocity outflow is mainly visible in lines requiring higher excitation. This
is however not a strict rule, as the high-velocity component is readily seen
in the doublet of Na I, an element that is highly ionized at higher
temperatures.

###  9.2 Reddening

In both lines of the sodium $D_{1}D_{2}$ doublet, we find an extra absorption
component centered at a heliocentric velocity of about –8 km s-1 that has an
FWHM of 13 km s-1 (at an instrumental FWHM of about 7.5 km s-1, its intrinsic
with is of 10.6 km s-1). This narrow component, although weak (with an
equivalent width, EW, of 0.034 Å), is also recognizable in the K I resonance
lines (see Fig. 17), but not in other lines. It may represent a cool
circumstellar or interstellar component. Given the Galactic position of U Equ,
the low radial velocity of this component supports its interstellar origin.
Indeed, an H I spectrum at 21 cm toward U Equ obtained by the Leiden-
Argentina-Bonn survey (Kalberla et al., 2005) shows a single feature at the
same velocity as the sodium lines. The depth of the component in both lines of
Na I is nearly the same, with EWs of 0.269 and 0.229 Å. Since the line ratio
in the optically thin case should be 2.0, the observed lines must be
saturated. The line of K I is, however, optically thin. Using the $E(B-V)$ vs.
EW calibration from Munari & Zwitter (1997) for Na I and K I, we obtain
$E(B-V)$ of 0.10$\pm$0.2 mag or $A_{V}\approx$0.3 mag. This extinction is
slightly higher than derived from extinction maps in Sect. 4.

We also attempted to constrain the total extinction by comparing the
intensities of the Balmer lines. The Balmer lines are too complex to
disentangle emission from absorption, and the entire normalized profiles were
measured for the comparison. The corresponding continuum levels (of
2.08$\pm$0.05 and 2.34$\pm$0.06 in units of 10-14 erg cm-2 s-1 A-1) were taken
from flux-calibrated Gaia spectra near the two lines. The H$\alpha$ to
H$\beta$ flux ratio we obtain, 3.0$\pm$0.1, is within the uncertainties
consistent with the Case B ratio of those lines at temperatures between 5
000–10 000 K. It is thus consistent with none or minimal extinction between us
and the recombining gas, consistent with our assumed $A_{V}$=0.24 mag.

###  9.3 Optical line variability

In principle, our spectra allow us to test for line variability over
timescales from days to about a year. Although spectra acquired a few days
apart show some differences in line intensities relative to the local
continuum, uncertainties in normalization to the continuum level and due to
the modest S/N are large and do not allow us to draw firm conclusions on the
variability. (Note that slit losses and other instrumental erratic effects are
irrelevant for the comparison of continuum-normalized spectra.) Over the year
covered by the observations, it would appear, however, that emission in the
strong recombination lines of H I and in the Ca II triplet is getting weaker
with time by a few tens of %. Changes are observed in the high-velocity wings
of the strongest absorption lines. The absorption at heliocentric velocities
between about –300 to –100 km s-1 is generally getting weaker. This is seen in
lines of ionized species, for example in the lines of Fe I near 5018, 6347,
and 6371 Å, in the hydrogen lines of the Paschen series, and in both of the O
I triplets near 7770 and 8445 Å; a similar change is seen also in the high-
velocity gas probed by the resonance absorption in the Na I doublet. It would
appear that the change is not related to excitation or ionization degree and
could be interpreted as geometrical thinning of the outflow along the line of
sight, for instance, due to expansion.

##  10 Discussion

###  10.1 Are the molecules gone? Not quite

The main result of this study is the realization that U Equ has undergone a
major transition and lost its unique molecular optical spectrum that drew
attention to the object in the first place. The disappearance of molecules
could be perhaps most naturally explained if intensified or hardened radiation
of the central star (or stars) completely dissociated the molecules in the
circumstellar environment. However, our observations at longer wavelengths,
obtained with ALMA and the Effelsberg telescope, seem to contradict this naive
interpretation.

The ALMA observations show that some cool molecular gas remains around U Equ.
Estimating its mass would require information on the (distance and) gas
temperature, which is not available at the moment. The “survival” of this cool
molecular gas with the simultaneous disappearance of optical molecular bands
may suggest that perhaps only the inner parts of the molecular envelope or
disk were dissociated. The weak rotational emission in CO 3–2 shows also that
the cool molecular outflow has a higher expansion velocity than assumed in the
literature for the wind of U Equ (based mainly on OH maser spectra). The half
of the full width of the CO line is of about 30 km s-1, consistent with normal
winds of AGB stars and OH/IR objects. Our Effelsberg observation of the water
maser lines show that the masers are still present tin the system, however
they display most bizarre variability. Circumstellar H2O masers are often
linked to shocks (e.g., Tafoya et al., 2014), but are known to be present in
systems with relatively hard spectrum, for instance, near high-mass
protostars.

The rather surprising discovery of an SiO maser ($\varv$=1) made with the
Effelsberg telescope complicates the picture even further, because, for an
evolved star, the masing action is expected to occur only a few stellar radii
from the stellar surface, as typically observed near M-type giants and
supergiants (see, for example Cotton et al., 2004; Reid & Menten, 2007).
Optical spectra definitely show a hotter photosphere in U Equ, of F type
(Sect. 9) or even hotter if the unconstrained SED models can be given credit
(Sect. 7). Any imaginable drastic deactivation of the molecular emission,
which is seen in the optical and at an excitation temperature of $\approx$500
K, would also likely switch off the SiO maser. These paradoxes can partially
be explained if U Equ is a binary and contains an M-type mass-losing giant
that has at least one order of magnitude lower luminosity than the F type
giant and thus does not show any photospheric features. Then, the observed
spectral changes would be related to clearing off the surroundings of the
hotter star, while leaving the envelope of the M giant unaffected. However, we
find such an interpretation rather unlikely, as an M giant with a wind thick
enough to produce an SiO maser (densities of 109 cm-3) would also be luminous
and should show some spectral features in the optical SALT spectra.

Alternatively, SiO may arise in a disk or disk-wind, like in the peculiar
young stellar object Source I in the Orion Kleinmann Low nebula (KL) (Matthews
et al., 2010; Issaoun et al., 2017). Therein, the central source is thought to
be a progenitor of a B-type star (Goddi et al., 2011; Ginsburg et al., 2018)
and may be characterized by a radiation field that is harsher than that of
M-type giants and closer to that of the main stellar component of U Equ. The
SiO masers are excited in an expanding and rotating disk-wind region of a size
of 40 AU with a possible involvement of circumstellar shocks. Because Source I
belongs to a disintegrated multiple system with its former members running
away from a place of a major eruption that took place $\approx$500 years ago,
it was proposed that Source I underwent a major dynamical interaction at this
time (Gómez et al., 2008). Some authors go as far as to propose that Source I
is a product of a merger (Farias & Tan, 2018). SiO masers like those in Source
I are however very rare (Zapata et al., 2009). Can U Equ be in some aspect
similar to this rare occurrence of circumstellar SiO masers? Has the
transition in U Equ triggered the maser emission only recently? What is the
relation of the SiO maser to the OH and H2O masers observed decades ago in U
Equ? Although a lot is unclear, it is apparent that U Equ is important for our
understanding of maser evolution in post-main sequence systems.

With U Equ having lost its molecular emission bands that sparked interest in
this object in the 1990s, there remain only a handful of objects with this
type of spectrum. Lloyd Evans (1997) discussed this small, but diverse group
of objects, which includes the RV Tauri stars and the red supergiant VY CMa.
However, in the 26 years since that discussion, the list of objects has grown
somewhat longer. VY CMa remains the poster child of this phenomenon, with
variable but persistent emission bands of TiO, AlO, ScO, and VO (Wallerstein &
Gonzalez, 2001; Kamiński et al., 2013). The Galactic red-nova remnants, V838
Mon, V4332 Sgr, V1309 Sco, products of stellar merger events, show similar
emission bands (e.g., Kamiński et al., 2015a), but those are fading on time
scales of decades. A few young stellar objects display emission bands after
outbursts (Herczeg & Hillenbrand, 2014; Hillenbrand et al., 2012), but they
are usually weak and fade away relatively fast. Mira stars also occasionally
show transient emission bands from material shocked by pulsations (e.g.,
Kamiński et al., 2016). Molecular emission bands are therefore a very rare and
– for most objects – a short-lived transient phenomenon. This was likely the
case of U Equ.

###  10.2 Nature of the photometric rise

The remarkable revamping of the optical spectrum associated with the seemingly
modest $\sim$1 mag change in the optical and near-IR fluxes could have been
caused by several phenomena. It appears to be possible that the change was
simply caused by a geometric effect, namely, a reorganization of the
circumstellar medium on the line of sight due to disk rotation or
inhomogeneous and anisotropic mass loss. In some scenarios, such a
reconfiguration would change the extinction and optical spectrum without any
major changes in the physical properties of the star. However, the
simultaneous rise at visual and mid-IR bands that we observe (with the caveat
that the behavior at longer wavelengths is unknown) could suggest an increase
in the object’s bolometric luminosity. Therefore, we rather favor a scenario
in which the transition was more fundamental than barely a change in
circumstellar arrangement along the line of sight. We find it more likely that
it is related to a major change in the circumstellar medium and to an increase
of the total luminosity of the system due to physical changes in the star or
due to activation of an accretion episode (on the main star or on its
hypothetical companion).

The amplitudes of the rise in photometric brightness in the visual and near-IR
bands that we find in U Equ are higher than amplitudes of typical long-term
flux changes in other post-AGB stars. Photometric monitoring projects
targeting several post-AGB stars show typical optical variations of
$\lesssim$0.5 mag (e.g., Mesler & Henson, 2008; Arkhipova et al., 2010;
Hrivnak et al., 2010, 2022). There are reports of changes that reach a higher
amplitude, but these are almost always short-lasting drops in brightness with
unremarkable changes in the optical spectrum. They can often be ascribed to
eclipses, trivial episodes of obscuration (for instance, by dusty wind
material), or to pulsations (cf. Şahin & Jeffery, 2007; Hrivnak et al., 2022).
Such cases do not require any activation of an extra energy source. The case
of U Equ is different because it has shown a continued, monotonic, long-
lasting rise by 1 mag associated with a dramatic change in the visual
spectrum. Such a high-amplitude monotonically rising brightness is certainly
not common among post-AGB objects, as it requires an extra energy sources.
Even if the relative change in the magnitudes scale is not huge, like here
$\approx$ 1, it implies a flux increase by a factor of a few, which at the
luminosity of $10^{4}$ L☉ implies a considerable amount of extra energy.

The possibility that we caught U Equ’s central star at the onset of the fast
transformation from an AGB object to a proto-WD is most fascinating. However,
with current data, we have no solid evidence of an increasing temperature at a
nearly constant luminosity of the star. The quality and coverage of the
available data is insufficient. However, we calculate that a black-body of a
fixed luminosity would increase its optical flux near 5000 Å with an
increasing temperature, so the observed photometric changes we observe are not
inconsistent with an early transformation on the post-AGB. (For instance, a
change from 4000 to 8000 K increases the flux by a factor of 2.3.)

The best observed fast transition in post-AGB stars is a manifestation of a
late He flash, where a cooling central star of a planetary nebula turns back
to a configuration with temperature and luminosity characteristic of a post-
AGB object. Several such ‘born-again’ (a.k.a. late-thermal-pulse) objects are
known (e.g., V4334 Sgr, V605 Aql, and FG Sge), and they all are quickly
evolving carbon stars (e.g., Van Winckel, 2003). U Equ is a descendant of an
O-rich OH/IR star (thus likely more massive than carbon stars) and the
characteristics of its photometric changes cannot be ascribed to a late He-
shell flash. A thermal pulse in an evolved massive AGB is another scenario
that can be considered for U Equ but is hard to verify with the data currently
in hand.

Like in many other post-AGB stars, we may be witnessing in U Equ an activation
of an accretion episode that increased the luminosity of the system, possibly
triggered by binarity and matter being accreted from a circumbinary disk. Such
systems are often associated with fast jets (Bollen et al., 2022). There is no
direct evidence of a binary companion to U Equ, but there hardly ever is for
post-AGB stars with dwarf companions of a much lower luminosity. Any
observational indication of jets in U Equ would confirm active accretion
episodes, and thus can explain some of the changes seen in this object over
the last decades. The short-term low-amplitude light variations observed in U
Equ by TESS (Sect. 3) is similar to flickering observed in accretion-active
sources (Scaringi et al., 2015), supporting our hypothesis of activated
accretion in U Equ.

Figure 7: Visual light curve of V576 Car extracted from ASAS and ASASSN data.
The rise in brightness has continued since at least 1989 (Couch et al., 2003).

###  10.3 U Equ as a post-AGB object

The general observational characteristics of U Equ, primarily its SED and its
spectral type, make it most similar to the class of post-AGB and post-RGB
objects known to harbor dusty disks. This class of objects has been recently
reviewed by Kluska et al. (2022). More specifically, U Equ resembles their
Category 0 sources, which are associated with disks seen at high inclinations.
If confirmed as a member, U Equ would be the first halo object in this group.
The post-AGB systems discussed by Kluska et al. (2022) show a chemical anomaly
whereby refractory metals are underabundant in the stellar atmosphere due to
selective accretion from a circumbinary disk. Future studies of U Equ should
investigate whether it displays a similar chemical peculiarity.

As noted in Sect. 9, we find a particularity strong resemblance of U Equ’s
contemporary spectra to those of the post-common-envelope binary system HD
101584. Multiple lines in the spectrum of HD 101584 are variable (Kipper,
2005) at a level similar to the changes observed in our SALT spectra. Spectral
features arising in the circumstellar medium of HD 101584 have multiple
components, with a maximum projected expansion velocity of 130 km s-1. This
velocity is only slightly lower than the extreme components of U Equ (215 km
s-1) and comparable to the outflow velocity measured for many strong
absorption lines (it is very likely that we look at the two systems from
different angles and projection effects are important). The central stars of
both systems may also be similar. Early UV spectra of HD 101584 seemed to
suggest a presence of a hot B9 II star (Bakker et al., 1996), but later
studies, based mainly on visual data, preferred a cooler star, of spectral
types A–F (Kipper, 2005), again in accord to what we know about U Equ. HD
101584, just like U Equ, has OH 1667 MHz masers.191919The most recent
detection observations are from te Lintel Hekkert et al. (1992). They are
located in a bipolar structure and display exceptionally broad line profiles
(te Lintel Hekkert et al., 1992). The projected OH expansion velocity of 42 km
s-1 is much smaller than the velocity of the CO outflow traced at (sub)mm
wavelengths (Olofsson et al., 2019) and the outflow probed by optical lines.
These unique masers bear strong resemblance to the unusual masers of U Equ
(except for the spatial distribution, which for U Equ is unknown; see Sect.
2.1). It is unclear whether HD 101584 has ever gone through spectral and
photometric transformations such as those reported in this paper for U Equ.
Since HD 101584 has been considered to be a rare example of a system that has
recently gone through a common-envelope phase (Olofsson et al., 2019), we
should consider U Equ as another unique system of similarly interesting recent
evolution and therefore worth of further studies.

###  10.4 V576 Car

The closest known analog to U Equ, as it was known to 1990s observers, is V576
Car (=IRAS 08182–6000). Its optical spectrum, dominated by molecular
absorption bands, but devoid of emission features seen earlier in U Equ was
described in detail in Couch et al. (2003). The circumstellar absorption bands
of all the main electronic systems of TiO, AlO, and VO were seen in the
spectra of V576 Car. The underlying stellar photosphere was not easy to
identify (as in U Equ), but the best efforts indicated an F–G supergiant,
similar to what we find for U Equ (Sect. 9). Additionally, Couch et al. (2003)
reported that the $J$-band fluxes were monotonically increasing over the
1990s, totaling to a 2.5 mag amplitude over 10 years. We found visual
photometric data displaying a continuation of this rise in brightness over the
last two decades by another 2.5 mag. The data taken from the All Sky Automated
Survey (ASAS; Pojmanski, 1997) and ASAS-SN sky patrol are presented in Fig. 7.
The steady rise in brightness seen in V576 Car is reminiscent of the light
curve of U Equ, although larger in amplitude. Additionally, TESS data for V576
Car (not shown) display the same small-amplitude erratic variability as that
we found in U Equ (Fig. 4). Most importantly, when we obtained a spectrum of
this object on 16 May 2023 with HRS-SALT, the absorption bands of molecules
had completely disappeared and been replaced by a spectrum remarkably similar
to the current spectrum of U Equ. The contemporary spectra of both objects are
shown along each other in Figs. 10–13. With little doubt, the changes in both
stars are very similar and in both cases resulted in an object with a spectrum
that is now more typical for post-AGB stars. The question that arises is how
common is the behavior when a post-AGB star switches from a spectrum dominated
by molecular bands into the atomic spectrum. The absence of masers around V576
Car (Couch et al., 2003) needs to be revisited with new sensitive
observations.

A larger sample of stars exhibiting the U Equ phenomenon would definitely help
to understand better the nature of the changes. Smith (2006) searched the
southern sky for objects with similar characteristics to U Equ and V576 Car
(i.e., as it appeared in the 1990s) but found none that would exhibit
circumstellar molecular features. The phase in which U Equ and V576 Car were
caught must therefore be short-lasting, as expected for fast-evolving dust-
embedded post-AGB stars. Perhaps future searches can focus on the steady rise
in brightness over timescales of $\gtrsim$3 decades as the identifying feature
of the U Equ phenomenon.

###  10.5 Summary and outlook

More observations are needed to explain the changes in U Equ and V576 Car.
Continued photometric and spectroscopic monitoring can reveal how long and to
what end the noticed changes will occur. Multiwavelength observations and
interferometric imaging would be particularly useful for understanding the
current configuration of the star or binary with respect to the disk and other
circumstellar features. We should be also looking for more objects
experiencing the U Equ phenomenon, as such studies have the potential to
expose the genuine and long-sought transition from AGB or OH/IR phase into a
planetary nebula, or they can deepen our understanding of common-envelope
interactions.

###### Acknowledgements.

T.K. acknowledges funding from grant SONATA BIS no. 2018/30/E/ST9/00398 from
the Polish National Science Center. K.I. was supported by Polish National
Science Center grant Sonatina 2021/40/C/ST9/00186. Based on observations
collected at the European Organization for Astronomical Research in the
Southern Hemisphere under ESO program 266.D-5655(A). This work made use of the
HDAP which was produced at Landessternwarte Heidelberg–Königstuhl under grant
no. 00.071.2005 of the Klaus-Tschira-Foundation. Some observations reported in
this paper were obtained with the Southern African Large Telescope (SALT).
Polish participation in SALT is funded by grant no. MNiSWDIR/WK/2016/07. Based
on data from the OMC Archive at CAB (INTA-CSIC), pre-processed by ISDC and
further processed by the OMC Team at CAB. The OMC Archive is part of the
Spanish Virtual Observatory project. Both are funded by
MCIN/AEI/10.13039/501100011033 through grants PID2020-112949GB-I00 and
PID2019-107061GB-C61, respectively. This research has made use of the SIMBAD
database, operated at CDS, Strasbourg, France. This research made use of
hips2fits, https://alasky.u-strasbg.fr/hips-image-services/hips2fits a service
provided by CDS. The CSS survey is funded by the National Aeronautics and
Space Administration under Grant No. NNG05GF22G issued through the Science
Mission Directorate Near-Earth Objects Observations Program. The CRTS survey
is supported by the U.S. National Science Foundation under grants AST-0909182
and AST-1313422. The Pan-STARRS1 Surveys (PS1) and the PS1 public science
archive have been made possible through contributions by the Institute for
Astronomy, the University of Hawaii, the Pan-STARRS Project Office, the Max-
Planck Society and its participating institutes, the Max Planck Institute for
Astronomy, Heidelberg and the Max Planck Institute for Extraterrestrial
Physics, Garching, The Johns Hopkins University, Durham University, the
University of Edinburgh, the Queen’s University Belfast, the Harvard-
Smithsonian Center for Astrophysics, the Las Cumbres Observatory Global
Telescope Network Incorporated, the National Central University of Taiwan, the
Space Telescope Science Institute, the National Aeronautics and Space
Administration under Grant No. NNX08AR22G issued through the Planetary Science
Division of the NASA Science Mission Directorate, the National Science
Foundation Grant No. AST-1238877, the University of Maryland, Eotvos Lorand
University (ELTE), the Los Alamos National Laboratory, and the Gordon and
Betty Moore Foundation. This research has made use of the VizieR catalogue
access tool, CDS, Strasbourg, France (DOI : 10.26093/cds/vizier). The original
description of the VizieR service was published in 2000, AAS 143, 23. Based 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 This research has made use of the
Spanish Virtual Observatory (https://svo.cab.inta-csic.es) project funded by
MCIN/AEI/10.13039/501100011033/ through grant PID2020-112949GB-I00.

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##  Appendix A Photographic plates

The field around U Equ was captured by photographic plates throughout the
first half of the 20th century. A sample of them is shown in Fig. 8.

Figure 8: Sample Heidelberg-Königstuhl blue plates. The cutouts are centered
near U Equ (encircled). The dates of observations are given. For reference, a
deeper POSS image is shown in the lower right corner. Images are displayed in
a linear intensity scale.

##  Appendix B Spectral energy distribution models

Figure 9: Model and observed SEDs of U Equ in 1989 (top), 2010 (middle), and
2023 (bottom). Observations are marked with points. Solid lines show best-fit
SED models from sedfitter for different circumstellar configurations descibed
in (Robitaille 2017b).

##  Appendix C Spectral atlas

Here we present an atlas of spectra of U Equ and V576 Car compared to that of
HD 101584.

Figure 10: Average normalized spectrum of U Equ (blue) is compared to that of
HD 101584 and of V576 Car (orange). Main telluric bands are marked in green.
The vertical dashed lines and corresponding labels identify the main spectral
features (”1” corresponds to neutral atoms and ”2” to first ionization state).
Figure 11: Continued. Figure 12: Continued. Figure 13: Continued.

##  Appendix D LTE analysis of Fe I lines

Using the ”Saha-LTE Spectrum” tool available from the NIST online database, we
simulated the intensities of Fe I lines at different temperatures between
4000–10 000 K with a step of 200 K. As the calculated line intensities are
measured in arbitrary units, the intensities were scaled to the relative
intensity of the line near 8514 Å, which has the highest S/N and is a blend of
two strong transitions. We analyzed only lines in pure emission, with no
blending or contamination from telluric or other features. The observed line
intensities were measured in IRAF. A $\chi^{2}$ test was run to find the best
fitting line ratios. The test yielded a temperature of 4200$\pm$200 K.

As a test of opacity, we also compared the line ratios of lines in the same
multiplet with the ratios of the transition probabilities, $A_{ki}$. If the Fe
I emission is mostly optically thin, the scaled fluxes and $A_{ki}$ values
should be comparable. The resulting comparison shows some agreement of the
NIST models within 5$\sigma$ errors, but not for all line pairs. The goodness
of fit shows no trend with upper energy level, but the majority of lines used
for this test originate from relatively close energy levels in range of 2–6
eV. Out of the nine Fe I ratios measured, four fit well within the 5$\sigma$
errors whilst three more do not lie much further from the observed values.
Only two line ratios vary significantly (6230/7748 and 6494/9593). We conclude
that the Fe I emission is not very optically thick.

##  Appendix E Line profiles

Figure 14: Profiles of selected lines. The dotted vertical lines mark the
center of the narrow absorption component discussed in the text. Figure 15: As
Fig. 14. Figure 16: As Fig. 14. Figure 17: As Fig. 14. Additionally, the
scaled and inverted profile of the H I emission line at 21 cm is shown with a
black line. The interstellar component is marked with a dashed line. Figure
18: As Fig. 14.