# Spectroscopic Signatures of the Vanishing Natural Coronagraph of eta Carinae

A. Damineli1, F. Navarete1, D. J. Hillier2, A. F. J. Moffat5, M. F.
Corcoran9,10, T. R. Gull3, N. D. Richardson4, G. Weigelt6, P. W. Morris7, I.
Stevens8,
1 Instituto de Astronomia, Geofísica e Ciências Atmosféricas da USP, Rua do
Matão 1226, Cidade Universitária, São Paulo, Brasil
2 Department of Physics and Astronomy & Pittsburgh Particle Physics,
Astrophysics, and Cosmology Center (PITT PACC),
University of Pittsburgh, 3941 O’Hara Street, Pittsburgh, PA 15260, USA
3 Laboratory for Extraterrestrial Planets and Stellar Astrophysics, Code 667,
NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA
4 Department of Physics and Astronomy, Embry-Riddle Aeronautical University,
3700 Willow Creek Road, Prescott, AZ 86301, USA
5 Département de Physique and Centre de Recherche en Astrophysique du Québec
(CRAQ)
Université de Montréal, C.P. 6128, Succ. Centre-Ville, Montréal, Québec, H3C
3J7, Canada
6 Max Planck Institute for Radio Astronomy, Auf dem Hügel 69, D-53121 Bonn,
Germany
7 California Institute of Technology, IPAC M/C 100-22, Pasadena, CA 91125, USA
8 School of Physics & Astronomy, Univ. of Birmingham, Edgbaston, Birmingham,
B15 2TT, UK
9 CRESST II & X-ray Astrophysics Laboratory, Code 662, NASA Goddard Space
Flight Center, Greenbelt, MD 20771, USA
10 Institute for Astrophysics and Computational Sciences, Department of
Physics
The Catholic University of America, Washington, DC 20064, USA E-mail:
<EMAIL_ADDRESS>
(May 10 2021)

###### Abstract

Eta Carinae is a massive interacting binary system shrouded in a complex
circumstellar environment whose evolution is the source of the long-term
brightening observed during the last 80 years. An occulter, acting as a
natural coronagraph, impacts observations from our perspective, but not from
most other directions. Other sight-lines are visible to us through studies of
the Homunculus reflection nebula. The coronagraph appears to be vanishing,
decreasing the extinction towards the central star, and causing the star’s
secular brightening. In contrast, the Homunculus remains at an almost constant
brightness. The coronagraph primarily suppresses the stellar continuum, to a
lesser extent the wind lines, and not the circumstellar emission lines. This
explains why the absolute values of equivalent widths (EWs) of the emission
lines in our direct view are larger than those seen in reflected by the
Homunculus, why the direct view absolute EWs are decreasing with time, and why
lower-excitation spectral wind lines formed at larger radii (e.g Fe ii 4585Å)
decrease in intensity at a faster pace than higher excitation lines that form
closer to the star (e.g. H$\delta$). Our main result is that the star, despite
its 10-fold brightening over two decades, is relatively stable. A vanishing
coronagraph that can explain both the large flux evolution and the much weaker
spectral evolution. This is contrary to suggestions that the long-term
variability is intrinsic to the primary star that is still recovering from the
Great Eruption with a decreasing mass-loss rate and a polar wind that is
evolving at a slower pace than at the equator.

###### keywords:

ISM: dust, extinction ISM: molecules – stars: binaries – stars: individual:
eta Carinae – stars: winds

##  1 Introduction

The evolution and fate of massive stars are often determined by binary
interaction (e.g. Sana et al., 2012). Their scarcity in the solar
neighbourhood makes progress in the field difficult. $\eta$ Car is relevant
because of its proximity (2.3 kpc), its long observational history, and
because it is one of the most massive stars within our Galaxy. It owes its
fame to the great eruption (GE) that occurred in the 1840s, when it became the
second brightest star in the night sky and ejected the elegant bipolar
Homunculus nebula (Davidson & Humphreys, 1997). Hidden behind circumstellar
dust (Andriesse et al., 1978; Viotti et al., 1989), the central object is a
highly-eccentric massive binary (Damineli, 1996; Damineli et al., 1997;
Damineli et al., 2000) with periastron passages occurring every 5.538 yr and
that are monitored throughout the electromagnetic spectrum (Pittard &
Corcoran, 2002; Corcoran et al., 2017; Duncan & White, 2003; Fernández-Lajús
et al., 2010; Mehner et al., 2014; Damineli et al., 2008; Teodoro et al.,
2016).

The primary star ($\eta$ Car A) dominates the optical light of the system. Its
extinction-corrected apparent magnitude would be $V$ = 0.94 mag, according to
modelling of the optical/UV spectrum (Hillier et al., 2006), plus a foreground
extinction of up to $A_{\rm V}$ = 2.5 mag (Teodoro et al., 2020). The apparent
magnitude of the central stellar core was $V$ = 7.73 on 1999.140 (Martin &
Koppelman, 2004), indicating additional local extinction obscuring our view.
If the dust was uniformly distributed around the stellar core, the Homunculus
reflection nebula would be proportionately fainter. The simple fact that the
Homunculus is so prominent, compared to the central star, indicates extra
extinction in our line of sight (LOS) to the central star. Moreover, the
relative high brightness of the three speckle objects BCD (Weigelt &
Ebersberger, 1986) could be explained if they are illuminated by the central
star and seen in an unimpeded view away from our LOS, while the central star,
located a fraction of an arcsec away, is seen under heavy extinction caused,
for example, by an obscuring dusty disk (Weigelt et al., 1995) or some other
optically thick, localized structure.

From the 1940s through the 1990s, the $\eta$ Car system (star + Homunculus)
had continuously brightened at a rate of $\delta V$ $\sim$ $-$ 0.02 mag yr-1
in the $V$-band. However, existing imagery from those epochs cannot separate
the stellar contribution from the nebula. Over the past two decades, the
brightening increased to a rate of $\delta V$ $\sim$ $-$0.05 mag year-1
(Martin et al., 2006), and improved photometric imagery allowed the separation
of the stellar core from the flux of the Homunculus nebula: the stellar core
brightened at a rate $\delta V$ $\sim$ $-$0.11 mag yr-1 while the Homunculus
nebula, has remained fairly constant at $V$ = 5.48$\,\pm$ 0.05 mag (Damineli
et al., 2019). The light emitted by the central star and reflected in the NW
extremity of the Homunculus reaches the Earth with a delay of $\sim$ 6 months.
If the star was intrinsically brightening, the Homunculus light-curve should
show the same effect delayed by a few months. In 2010 the brightness of the
central star surpassed that of the Homunculus in the $V$-band, for the first
time in 160 years. Since the Homunculus is a reflection nebula and is almost
constant in brightness, the extinction along our LOS to the central star must
be changing (Damineli et al., 2019). The brightening cannot be explained by
the central star becoming intrinsically more luminous since the stellar
spectrum has not evolved significantly even though the object brightened by
approximately a factor of 10 in the last 70 years. No such long-term
brightening was observed in the mid-infrared, as reported by Mehner et al.
(2019). The sparse observations in the mid-infrared since 1968 are compatible
with lower amplitude variations (factor $\sim$ 2) on a timescale of years, as
is the case also for the orbital modulation in the optical reported by
Damineli et al. (2019). The mid-infrared flux is produced by the absorption
and re-emission of light from the central object by dust spread over the whole
nebula. In contrast to the optical and ultraviolet emission, it is extended,
and not affected by any localised structure.

$\eta$ Car is a strong X-ray source, in which thermal X-rays are produced in
the hot shocked gas where the wind from the primary star collides with the
wind of the companion. The luminosity and temperature of the X-ray emission
depend on the wind properties (i.e. mass-loss rates and pre-shock wind
velocities: see for example Stevens et al. (1992). The wind parameters in turn
depend on the underlying stellar parameters. Thus, variations in the intrinsic
stellar properties would be reflected in variations in the observed X-ray
spectrum. The X-ray temperature and luminosity near apastron (when Coriolis
distortions are minimal) show little variations from orbit-to-orbit over the
last four orbits for which detailed monitoring data exist - see Corcoran et
al. (2017) and Espinoza Galeas et al. (2021) \- indicating that neither the
wind nor stellar parameters have varied significantly during the prior two
decades. This conclusion is supported by the lack of variations in the
terminal velocities derived from the P-Cygni absorption profiles of the upper
members of the Balmer series (like H$\delta$) (Teodoro et al., 2012).

The initial suggestion that we see the central star obscured by dust
condensation and that the Homunculus scatters radiation that has suffered
relatively little attenuation was made by Hillier & Allen (1992). This idea
arose because the EWs of the lines seen in direct light from the central star
were larger than the corresponding EWs measured in reflection from the
Homunculus. This was true for some wind lines (Fe ii, H i), and particularly
true for the narrow nebular lines. Indeed, the reflected spectrum was much
more similar to a classic P Cygni spectrum than the spectrum of $\eta$ Car
observed in direct light. Hillier et al. (2006) discussed additional evidence
of the presence of a kind of natural coronagraph in $\eta$ Car.

Based on spatially-resolved HST/STIS spectroscopy of the star, Hillier et al.
(2006) showed that the spectrum of $\eta$ Car could be modelled only if the
star and its wind suffered a visual extinction ($\approx 7$ mags at V) that
was much larger than could be attributed to the interstellar medium ($<2$
mags). However, the models of Hillier et al. (2006) always underestimate the
observed EWs of some emission lines such as H$\alpha$ and H$\beta$. The
inability to match these lines can be attributed to the presence of an
occulter that reduces the continuum flux more than the line flux.

The presence of the coronagraph explains the unusual spectra taken in the
1980s and 1990s using ground-based spectrographs. The spectra exhibited broad
hydrogen and permitted iron lines characteristic of a typical P Cygni star
with a terminal velocity of $\sim$ 500 km s-1 (e.g. Zanella et al., 1984;
Hillier & Allen, 1992). However, the spectrum also showed narrow forbidden
lines of Fe superimposed on broad profiles. The narrow forbidden lines arise
in partially ionised structures including the three Weigelt knots (Weigelt &
Ebersberger, 1986; Davidson et al., 1995), while the broad components arise in
the very extended wind (Hillier et al., 2001). Normally the wind components
would be too faint to be seen. However, the coronagraph systematically weakens
the stellar continuum, but less so the broad components from the very extended
wind. As the coronagraph vanishes, the stellar continuum emission increases
and hence the extended broadened wind lines and the narrow forbidden lines,
which have constant flux, become less prominent (Mehner et al., 2010b;
Davidson et al., 2005).

Mehner et al. (2012) made detailed monitoring of a set of spectral lines in
the period 1998-2011, in direct light to the star and also at the FOS4
position (Davidson et al., 1995) near the SE pole of the Homunculus. In
addition to confirming the weakening of the wind lines of the central star,
that work showed that the broad-line absolute EWs formed in the primary’s wind
decreased at a faster pace than those observed in reflected light at FOS4.
Part of the present work is focused on the equivalent width (EW) evolution of
the three representative lines of the stellar wind (H$\alpha$, H$\delta$ and
Fe ii 4585Å) studied by Mehner et al. (2012), extending the data set to the
time interval 1990-2021 (see Figure 4). The observational results of Mehner et
al. (2012), made during a shorter time interval, were confirmed, although the
physical mechanisms suggested to explain the observations differ from those
presented here.

Mehner et al. (2012) suggested that the secular evolution of the spectra is
due to variations intrinsic to the primary star that is still recovering from
the 1840s Great Eruption. They suggest that the mass-loss rate is decreasing,
making the wind more transparent. However, since the evolution of the spectral
lines from the star and the corresponding ones at FOS4 are different, Mehner
et al. (2012) postulate that the stellar wind at the equator is evolving
differently from that at the pole. While the P Cygni absorption component can
be strongly dependent on viewing angle, the emission component in a moderately
aspherical wind will vary much less, simply because one is averaging the
emission over a large volume. Thus, it is unlikely that a simple latitudinal
variation in mass-loss rate could explain the observed variations. Further,
this model, by itself, does not explain the large changes in the apparent
brightness of $\eta$ Car.

Recent ALMA observations have also provided direct confirmation of the
coronograph at sub-arcsecond resolutions ($<$ 0$\aas@@fstack{\prime\prime}$5
$\sim$ 1100 au). Although ALMA mapped molecular absorption lines, they should
trace the warm/cold dust. Bordiu & Rizzo (2019) found CO ($-$9 km s-1), HCN
(Hydrogen cyanide) and H13CN ($-$60 km s-1) in absorption in front of the star
at $T$ $\sim$ 500 K. Morris et al. (2020) found an absorption spot they
attributed to CH3OH ((Bordiu & Rizzo, 2019)) and another of 12CO at radial
velocity $+$3 $\pm$ 5 km s-1 in front of the star at $T$ $\sim$ 100-110 K.
Different velocities and temperatures suggest that the absorbing material
(inside the occulter) are spread at different radii in front of the star.
These molecular absorption lines should be fading, but only future visits by
ALMA can test this hypothesis.

In this work, we present new data and analyse spectroscopic evidence for the
existence of a natural coronagraph in front of $\eta$ Car. The paper is
organised as follows: Sect. 2 presents the data; Sect. 3 presents our results;
in Sect. 4 we discuss these results; Sect. 5 presents our conclusions.

Through this paper we use the binary ephemerides reported by Teodoro et al.
(2016): period $=$ 2022.7d, phase zero $=$ HJD 2456874.4. Following Grant et
al. (2020) the actual periastron passage occurs four days earlier than this
definition of phase zero. The numbering of the orbital cycles was defined by
Groh & Damineli (2004).

##  2 Observations

We have collected a large number of medium/high-resolution spectra over the
past 30 years from different observatories, by different observers and
reported in a number of papers.

The reader should bear in mind that we do not aim to present a general
description of the spectral features in $\eta$ Car – this has been done in
earlier works. We selected the data by focusing on a specific question: Is
there any strong support for the existence of a natural coronagraph in $\eta$
Car? Consequently, we do not explore spectral features and particular modeling
results that are interesting but a distraction from our goal.

Homogeneity of the data is a major concern in a large data set. Also of major
concern when measuring EWs and line fluxes is the definition of the stellar
continuum. This is particularly difficult in $\eta$ Car since the lines are
relatively broad, frequently display extended wings, and are often blended.
Échelle observations can exacerbate the problem further because of the
difficulty in joining orders. Signal-to-noise ratio (S/N) and spectral
resolving power (R) are important for measuring line profiles, less so for
measuring EWs. For measuring EWs, the wavelength integration limits are
important, although subject to a somewhat arbitrary choice, due to blends with
neighbour lines. This can be controlled by always using the same limits. The
blended lines are much fainter and insert a low-level bias in the time series.
The seeing also is not too critical for the central star in $\eta$ Car, since
most of the emitting structures are confined to the central
0$\aas@@fstack{\prime\prime}$5, and the nebular contribution is far smaller.

As a quality check on measurements, we have compared slit spectra taken at
1′′-1$\aas@@fstack{\prime\prime}$5 seeing with FEROS fiber-fed (diameter
2$\aas@@fstack{\prime\prime}$7) and found no important line profile
differences. We made extractions of the stellar spectrum in a typical UVES
spectrum (seeing fwhm $\sim$ 0$\aas@@fstack{\prime\prime}$8) using apertures
starting from 0$\aas@@fstack{\prime\prime}$5 (2 pixels) up to 2′′ (8 pixels)
and found no significant differences in the normalised spectra. We then
defined 1$\aas@@fstack{\prime\prime}$0 as the length along the slit to extract
the UVES spectra. In order to check our procedure, we compared our
measurements with those published in Mehner et al. (2012) and found reasonable
agreement. Spectra from the star taken with different ground-based
spectrographs coeval with UVES observations also showed good agreement without
applying any adjustment shifts in EWs.

Good image quality is of major importance for the nebular emission, as is the
scattered light internally to the spectrograph. For this task, the UVES data
set taken at the FOS4 position is crucial.

HST/STIS spectra resolve the structures around the central star, which demands
special care in extracting the spectra to be compared with ground-based
observations. For the sake of homogeneity, since they do not exhibit a dense
time sampling or long-term monitoring as the ground-based data, we did not use
the HST data.

Figure 1:  Profiles of representative primary wind lines in direct (black) and
reflected (magenta) views. Spectra taken at high excitation phase
($\phi\,=\,$11.52, May/June 2006) with UVES. Panel a) H$\alpha$; b) H$\delta$;
c) Fe ii $\lambda$4585Å; d) High members of the Balmer series. Radial velocity
shifts have been applied in order to match the centroid positions of the
emission lines in the reflected spectrum with those in direct light.

For the purpose of the present work, the main sources of data are described
below.

  * •

UVES direct and reflected light spectra: We selected high R and S/N spectra
observed at VLT/UVES/ESO-Chile with $R$ $\sim$ 100,000. The data were
retrieved from the $\eta$ Car Treasury Program111http://etacar.umn.edu/
archive and were used to study the difference between direct and reflected-
light spectra. We extracted UVES 1D spectra by tracing the 2D spectra within
$\pm$2 pixels (1′′) from the slit center. The slit width was
0$\aas@@fstack{\prime\prime}$4\. We used coeval stellar and FOS4 spectra (when
available) with the same parameters to extract the nebular spectra. Most of
the UVES spectra at FOS4 (Davidson et al., 1995; Mehner et al., 2012) in the
period 1999.97-2009.50 were centered at $\sim$ 2$\aas@@fstack{\prime\prime}$84
E and $\sim$ $-$2$\aas@@fstack{\prime\prime}$36 S of the central star
(3$\aas@@fstack{\prime\prime}$7 SE). Unfortunately, after 2013 this position
was missed by the telescope pointing and the slit center was not in
coincidence with FOS4. As shown in Mehner et al. (2012), line profiles change
from point to point. Although extracted the few nebular spectra in 2013-2015
at $>$ 0$\aas@@fstack{\prime\prime}$5 off the slit center, we did not use
these in the formal analysis. EWs from FOS4 spectra are listed in Table 8.

A pair of spectra in direct and reflected light was taken in May/June 2006,
separated by 18 days, both during the high-excitation phase - when the
cyclical variability arising from the 5.538 yr period was at a minimum - to
extract the representative line profiles (they are shown in Fig. 1). The
observations of the direct spectrum were obtained on HJD = 2453894.5, and for
the reflected spectrum at the FOS4 position on HJD = 2453866.5. Although not
simultaneous with the reflected spectra, the inspection of H$\alpha$ spectra
taken in direct light at several observatories in the time frame 2005.5-2006.5
confirmed that H$\alpha$-EW was stable at better than 5%.

  * •

Optical spectra taken with the FEROS fiber-fed spectrograph ($R$ $\sim$
48,000) at the 0.52 and 2.2-m telescopes at La Silla Observatory (ESO-Ch): The
spectra were obtained in the time-frame 1992-1999 (Damineli et al., 2000, and
references therein) and a few others in subsequent years. In these spectra, we
measured EWs for H$\alpha$, H$\delta$ and Fe ii 4585Å. The fiber diameter
projected onto the sky was 2$\aas@@fstack{\prime\prime}$7.

  * •

Optical spectra taken with the CTIO 1.5m telescope with two different
spectrographs: CHIRON/CTIO ($R$ $\sim$ 90,000) in the period 2009-2021 and the
previous version of this spectrograph, a fiber-fed echelle with $R$ $\sim$
40,000; as described by Richardson et al. (2010); Richardson et al. (2015);
Richardson et al. (2016); Mehner et al. (2012); Teodoro et al. (2016).

  * •

The SOAR/Goodman spectrograph was used in several runs from 2008 to 2014 in
the blue part of the spectrum, with $R$ $\sim$ 3500 (Teodoro et al., 2016).

  * •

Optical spectra collected with NRES/LCOGT at South Africa and CTIO with $R$
$\sim$ 48,000. The fiber diameter projected onto the sky was
2$\aas@@fstack{\prime\prime}$8 (Navarete et al., in prep.). The 2019.9 LCO
spectrum was used to compare with the 1999.1 FEROS spectrum listed in columns
5 and 6 of Table 2.

  * •

Optical spectra were taken at the Coudé focus of the 1.6 m telescope at the
Pico dos Dias Observatory (OPD/LNA-Br) In the optical range a variety of CCDs
were used, covering $\sim$ 500 to $\sim$ 1000Å at each grating angle. The
response of the system as a function of wavelength is very flat and a simple
linear fit can be used to rectify the stellar continuum in each spectral
range, in contrast to the blaze function of the Echelle spectrographs. The
room temperature of the spectrograph is very stable, resulting in accurate
radial velocities. In the period 1989-2005 we used thick CCDs, delivering
R$\sim$6,000. After the year 2005 we used 2048 $\times$ 2048 thinned CCDs
delivering R$\sim$ 12,000. The typical slit width was $\sim$
1$\aas@@fstack{\prime\prime}$5 and the seeing varied from $\sim$
1$\aas@@fstack{\prime\prime}$0 to $\sim$ 3$\aas@@fstack{\prime\prime}$0,
although a number of observations were done at higher airmasses, which implies
somewhat worse seeing.

  * •

NIR spectra were taken at the Coudé focus of the 1.6 m telescope at the Pico
Dos Dias Observatory (OPD/LNA-Br) to survey the He i 10830Å line since 1989.
Before 2001 we used thick CCDs delivering $R$ $\sim$ 12,000, covering the
spectral range $\lambda$10700-11000 Å. Although having efficiency lower than
1%, these arrays did not exhibit noticeable fringes like the thinned CCDs,
which were used after 2005. The thinned CCDs delivered $R$ $\sim$ 22,000 and
covered approximately the same spectral range as the thick ones. In the period
2001-2004 a HgCdTe infrared array was used, delivering $R$ $\sim$ 7,000 (Groh
et al., 2007), covering the spectral range $\lambda$10450-11000 Å.

  * •

A NIR spectrum taken at CRIRES/VLT/ESO-Ch on 2009 April 03 and reported by
Groh et al. (2007) was used to explore the spectral absorption feature at
$\lambda$ 10792Å. The resolving power was $R$ $=$ 90,000 and spatial
resolution 0$\aas@@fstack{\prime\prime}$33).

##  3 Results

###  3.1 Spatially resolved line strengths

We checked the finding by Hillier & Allen (1992) who reported that spectral
lines scattered by the Homunculus suffer less extinction than those observed
in direct light towards the central star. UVES observations were carried out
with sub-arcsec spatial resolution, not requiring complicated correction for
the light scattered inside the spectrograph as in the previous work. Figure 1
shows representative lines we used.

A basic characteristic of the wind lines is that the higher the upper energy
level, the closer the line originates to the primary’s stellar core. Lines of
lower excitation arise from increasingly larger regions of the primary’s wind
(which dominates the wind emission), while higher excitation wind lines
originate in the inner wind near the star. Thus, the comparison of continuum-
normalised wind line-profiles of different excitation towards the star and
reflected by the Homunculus can be used to map the spatial extent of the wind.

Panels a through c in Fig. 1 display Balmer lines by increasing upper level.
The similarity between the direct and reflected light improves from lower-
level transitions (e.g. H$\alpha$ and H$\delta$, Fig. 1.a and Fig. 1.b)
towards upper-level transitions (Fig. 1.c). We use here model parameters for
the central star using CMFGEN, quoted in Sect. 4.1. The H$\alpha$ line (Fig.
1.a) forms up to large radii in the stellar wind ($r$ $\sim$ 100 R∗ $=$
12,000$R_{\odot}$ $=$ 56 au $=$ 24 mas). Therefore its wind component is
mostly insensitive to LOS absorption, contrary to the continuum flux, which
increases towards the star. Thus the apparent decrease in the continuum-
normalised strength of the H$\alpha$ line in the reflected spectrum (Fig. 1.a)
shows that the continuum suffers higher extinction in the LOS than in other
directions. On the other hand, the H$\delta$ line arises from a smaller region
around the star and therefore is sensitive to the same small-scale LOS
absorption that also affects the stellar continuum. Thus the temporal
continuum-normalised H$\delta$ line in the direct spectrum is more similar to
that reflected at FOS4 than the case of H$\alpha$ (Fig. 1.a).

In direct light, the upper-level transitions of the Balmer series (Fig. 1.d)
are very similar to those observed in the reflected spectrum, which indicates
that their formation region is entirely obscured by the coronagraph. The
transition H 12$-$2 at the right of Fig. 1.d is formed at $r$ $\sim$ 10 R∗,
which poses a minimum radius for the coronagraph, arbitrarily assuming it to
be circular. Figure 1.c shows the line profiles of permitted and forbidden
singly-ionised iron around $\lambda$4585Å. They are typically formed at $r$
$\sim$ 200 R∗, and suffer less extinction by the coronagraph than the stellar
continuum, similar to H$\alpha$. An extreme example of the contrast effect is
the fact that the He ii 4686Å EW is repeatable from cycle to cycle, as
reported by Navarete et al. (2020) and by Teodoro et al. (2016). At
periastron, the bulk of this Helium line is formed in a small region around
the wind-wind collision (WWC) apex, unresolved in HST/STIS spectra. Like the
upper-level transitions of the Balmer series, its characteristic disk is
smaller than the coronagraph, since He ii 4686Å EW does not show the varying
contrast effect between the stellar continuum and line emission. The He i
7065Å line displayed in panel e behaves differently from the others we have
described. Its P Cygni absorption is stronger in direct than in reflected
light because it has a contribution from the WWC walls (Richardson et al.,
2016).

Figure 2:  Variation of Fe ii $\lambda$5169 and [Fe ii] $\lambda$5159 narrow
and broad components at the same phase for four orbital cycles.
$\phi\,=\,$9.004 corresponds to 1992 June 20. a) Line profiles normalised to
the stellar continuum showing the long-term decreasing strength caused by the
secular brightening of the stellar continuum. b) Line profile fluxes relative
to the $\phi\,=\,$9.004 brightness ($V$ = 5.64 mag). The narrow line
components remained constant because they are formed in the Weigelt knots,
outside the coronagraph. c) Flux-normalised line profiles. The broad line
components, formed in the wind, are little affected by the brightening.

The strong, narrow emission lines superimposed on the direct stellar spectra
are produced by the Weigelt knots and other fainter dense structures spatially
resolved but in close vicinity to $\eta$ Car (Davidson et al., 1995; Gull et
al., 2016). They also appear in the spectra scattered by dust in the
Homunculus, but are greatly attenuated and significantly broadened as they are
intrinsically very faint. Their greatly decreased contribution in reflected
light is another indication that the stellar continuum along our LOS to the
primary core is largely suppressed.

Some features shown in these spectral lines may not be related to the
depression of the stellar continuum, and will not be further discussed. For
example, the distortion on the line profiles, like the P Cygni absorption and
the shoulder in the red arm of H$\alpha$ $\lambda$6563 Å and Fe ii lines, is
likely due to a latitude-dependent stellar wind (Smith et al., 2003c), or the
effects of the wind cavity (Groh et al., 2012c), and inherent broadening due
to velocity dispersion of the scattering dust.

###  3.2 Cycle-to-cycle weakening of the EW of the primary’s wind lines

We demonstrate in Fig. 2 how the continuum level has changed relative to line
profiles with two typical, singly-ionised iron lines $\lambda\lambda$ 5159 and
5169Å. Note that the first transition is permitted and the second is
forbidden. They have been observed at nearly the same phase over four orbital
cycles. This figure shows line profiles observed at 9 days after phase zero,
with phase $\phi\,=\,$9.004 corresponding to 1992 June 20. Figure 2.a shows
the rectified spectra for a region containing both a permitted and a forbidden
transition. Both lines exhibit a broad component formed in the wind, and a
narrow component that is formed in the BCD Weigelt knots located at $>$
0$\aas@@fstack{\prime\prime}$25 to the NE of the stellar core. As shown in
this plot, the line’s absolute EWs have decreased with time.

Figure 2.b shows these two lines’ flux corrected using the $V$-band coeval
magnitudes of the Homunculus (in the upper right) and scaled to unity at phase
$\phi\,=\,$9.004 when the brightness was $V$ = 5.64 mag. The flux of the
narrow-line components remained constant, confirming that the Weigelt knots
remained at a constant brightness. This demonstrates the complementarity of
ground-based observations to STIS/HST resolved spectra, as the space-based
observations were the first to reveal the roughly constant brightness of the
Weigelt knots (Hillier et al., 2006; Gull et al., 2009; Mehner et al., 2014).

Figure 2.c shows flux-normalised line profiles, which is the same as Figure
2.b when it is normalised to unity at the stellar continuum. This plot shows
that the broad line components formed in the wind are little affected by the
brightening. The broad lines have changed only slightly with time and, in
particular, there is no indication for a change in the speed of $\eta$ Car’s
wind. Since the Weigelt knots were not affected by the changes in the
continuum flux, the narrow line components in Figure 2.c are off-scale.

The natural explanation for these features is that the major change with time
is caused by decreasing extinction due to dissipation of dust and/or
tangential motion of the obscuring material covering the stellar core and
possibly the extended structure located to the SE, opposite the positions of
Weigelt C and D relative to $\eta$ Car.

###  3.3 Differential weakening of Balmer lines for the central core

As we have extensively monitored the star, especially across periastron
passages since 1992 (periastron 9), we can compare the evolution of line
profiles at the same binary phases. Spectra recorded at phase $\phi$ $\approx$
0.025 ($\sim$50 days after phase zero), have been taken when the secondary
star is deeply embedded in the extended primary wind as it passes behind the
primary star. Thus, we are able to probe the relatively undisturbed wind of
the primary star for a brief time (days to a few weeks) across each periastron
event.

Figure 3.a shows the evolution of the normalised H$\alpha$ line-profile across
six periastron events. The strongly continuum-emitting core has become much
more visible, causing the absolute EW of H$\alpha$ to decrease. By contrast
the H$\delta$ line shows only minor changes in the emission profile (Fig.
3.b), indicating that it is formed primarily behind the coronagraph.
Furthermore, the blueshifted P Cygni absorption, which is produced along our
LOS to the star, also has not changed. The constancy of the profile indicates
that the terminal velocity, the mass loss, and the luminosity (by inference)
of the star, have not changed very much over the last 30 years.

Figure 3:  Evolution of H$\alpha$ and H$\delta$ at periastron. Line profiles
of H$\alpha$ (panel a) and H$\delta$ (panel b) during the last six orbital
cycles at the same phase $\phi$ $\approx$ 0.025 ($\approx$ 50 days after
periastron). The H$\alpha$ line peak decreased by $\approx$ 50% while
H$\delta$ remained nearly constant. The constancy of the H$\delta$ P Cygni
absorption depth, terminal velocity and emission indicate a stable mass-loss
rate. There are no H$\delta$ observations at this phase for cycles #10 and
#13. The same colours in both panels correspond the same observing dates.

###  3.4 Evolution of line strengths in direct and reflected light

Figure 4 shows the time evolution of three representative lines of the
primary’s wind from 1990 to 2021: Fe ii 4585 Å in the top panel, H$\alpha$ in
the middle and H$\delta$ at the bottom. The EW absolute value222the absolute
value is used because the classical definition of EW is negative for emission
lines has decreased in some cases and remained constant in others. To aid in
the interpretation of the variations we fitted the data (using least squares)
to a simple linear function of the form $EW=a(t-2000)+b$, and the coefficients
of the fits are provided in Table 1.

Tables 5, 6 and 7 show all measurements for the star from all observatories.
EWs were measured for the entire set of spectra using the same parameters as
those defined by Mehner et al. (2012) for the continuum normalisation and EW
integration. The integration range was $\lambda$ 6520-6620 Å for H$\alpha$,
$\lambda$ 4085-4115 for H$\delta$ and $\lambda$ 4570-4600 Å for the Fe ii
$\lambda$4585 Å line.

To place lines of different strengths on a more equal footing we use the EW
ratio between the years 1995 and 2020 (R95/20 in column 4 of Table 1) to
compare variations. The Fe ii 4585 Å absolute EW decreases at a pace $\sim$ 6
times faster than H$\alpha$ and $\sim$ 11 times faster than H$\delta$. The far
smaller changes in H$\delta$ EW (which is formed close to the star) compared
to that of H$\alpha$ or Fe ii 4585 Å (formed far in the wind) indicates that
H$\delta$ is being emitted in a region that is mostly covered by the
coronagraph. It also suggests that across the continuum and H$\delta$
formation regions the occulter is fairly uniform. However, it does not need to
be concentric to the star.

Table 1:  Absolute EW rate of decrease compared to the formation radius for
spectral lines in direct view to the star

Feature | slopea | intercepta | R95/20b | log(r/R*)c
---|---|---|---|---
Units | Å.yr-1 | Å |  | -
H$\delta$ | 0.12 | -27.5 | 1.1 | 1.5
H$\alpha$ | 21.13 | -990.4 | 1.9 | 1.9
Fe ii 4585 | 0.581 | -12.9 | 12.3 | 2.3

Notes: a linear fit EW (Å) = a . (year - 2000) + b
b R95/20 = EW(1995) / EW(2020) using the linear fit
c CMFGEN model described in Sect. 4.1

Table 8 contains EWs of the same three lines reflected at FOS4. At this
position, the decreasing rate was $-$0.2 Å yr-1 for the absolute EW of Fe ii
4585 Å; the H$\alpha$ line strength remained constant at $-$523 Å; and
H$\delta$ also remained constant at $-$20 Å. The blue triangles at the right
end of the FOS4 series were extracted from spectra centered at distances
larger than 0$\aas@@fstack{\prime\prime}$5 from FOS4. Although they follow the
series reasonably well, were not used in the calculations.

Changes in the H$\alpha$ EW measured in direct light seem to have stalled
during the last orbital cycle (after 2014.5), with a value approaching the EW
of the Homunculus-reflected light. This is expected when the coronagraph has
completely vanished. But this is earlier than predicted by Damineli et al.
(2019) and may be due to a temporary halt in the dissipation/shifting of the
coronagraph.

Figure 4:  Cycle-to-cycle evolution of the equivalent width of representative
primary-star wind lines in direct (red) and reflected view at FOS4 (blue). Top
panel) Fe ii 4585 Å has the steepest gradient of growing fainter in both
views. Middle panel) H$\alpha$ direct shows an intermediate decrease rate of
the absolute EW and constant level $\sim$ $-$523 Åat FOS4. Bottom panel)
H$\delta$ in direct view shows a very gentle slope and at FOS4 the line is
almost constant at EW $\sim$ $-$20 Å. Gray points indicate measurements
reported by Mehner et al. (2012), with $+$ indicating direct view and $\times$
reflected at FOS4. Measurements in the range of phases 0.93-1.07 were omitted
in order to highlight cycle-to-cycle variations. Blue triangles indicate EWs
for recent years when the telescope pointing was too far from FOS4; these were
not used in the linear regression fits.

###  3.5 The blue-displaced absorption (shell-like feature) in the H$\alpha$
line profile

Figure 5:  Historical EW time series of the blue displaced absorption on the
H$\alpha$ line profile - originated in the LH - showing three components: a)
the periodic peak at periastron, b) occasional excursions to high intensity
after periastron, and c) a relatively low-intensity component always present,
but fading in the long term. Labels $\\#$09 to $\\#$14 indicate the orbital
cycles. The EW was measured by direct integration over the H$\alpha$ line
profile, in the range $\lambda$ 6548-6556 Å.

The blue absorption feature in the H$\alpha$ line profile (see Fig. 3.a) is a
signature of material in our LOS, which might be related to the coronagraph.
It was first reported by Ruiz et al. (1984) and detected periodically in all
subsequent low excitation events that occur near inferior conjunction of the
primary (Damineli et al., 1998; Richardson et al., 2010; Richardson et al.,
2015). It is not intrinsic to the primary wind; however, its velocity is
consistent with that of the Little Homunculus (LH) nebula (Ishibashi et al.,
2003; Gull et al., 2006). When the system is outside periastron, the LH
outside the LOS is fully ionised by the FUV emitted by the secondary star and
the WWC and escaping through the shock hole. However, the ionised hydrogen
recombines when the dense wind of the primary inhibits the escape of the
photons from the secondary, but allows Lyman $\alpha$ to scatter in the
circumstellar medium (Johansson et al., 2005). This occurs when there is a
neutral or partly ionised intervening gas in our LOS, which produces a typical
shell-like absorption, similar to that observed in some Be stars. The shell-
like feature which has RV $=$ $-$144 km s-1, FWHM $=$ 25 km s-1, was replaced
by a broader component in cycle #14, with FWHM $=$ 82 km s-1 centered at RV
$=$ $-$180 km s-1, which continues to be compatible with the LH velocity. This
indicates that the primary’s wind, or the region of the LH causing the shell-
absorption, suffered a change during the last periastron.

Figure 5 presents the EW historic time series of the shell-like feature
associated with H$\alpha$. The dashed lines indicate phase zero of the
periodic events. In addition to the periodic component, there are high-
intensity excursions of the EW after phase zero of cycles #09, #10 and #11.
These mid-cycle peaks have duration of four years or less and might be due to
dense structures passing across our LOS to the secondary star. They appear to
have not been present in cycle #13. The third feature in this time series is a
low-level absorption component which was present outside periastron and has
been decreasing in EW with time. It remained at EW $\sim$ 0 after the low
excitation event #13.

The periodic peaks of the blue H$\alpha$ absorption EWs (Figure 5) are not
produced by the coronagraph, which varies on longer timescales. The occulter
that causes the periodic peaks around periastron is the primary’s wind,
produced when the secondary star moves behind the primary near periastron. The
shell-like absorption is also observed outside periastron and has a long-term
behaviour. It is probably produced by other regions around the binary system.
The long-term variable shell-like absorption is produced by the coronagraph.
The complex behaviour of the shell-like absorption indicates that there are
other occulters located in the circumstellar medium, located farther from the
primary star but closer than the LH.

A similar narrow absorption line also exists in other directions, as reported
by Boumis et al. (1998). That direction could cross denser circumstellar
material inside the LH that is sampled by our LOS and must be permanently
neutral or partly ionised. This means that the UV radiation emitted through
the WWC cavity never illuminates that direction. That shell-like absorption is
reflected on a small nebula at $>$2′ to the NW of $\eta$ Car, outside the
Homunculus, and is seated at RV $=$ $-$80 km s-1 on the H$\alpha$ line
profile. It is stronger than its sibling along our LOS and has not varied
during 1985-1997. We took one spectrum with the Goodman/SOAR spectrograph in
2018 at the same spot as observed by Boumis et al. (1998). Although our
spectra have lower spectral resolution, no major changes were seen. This seems
to indicate that the shell-like feature sampled in our LOS is part of a much
larger structure.

Pickett et al. (2021) reported a similar decrease of the EW weakening of the
$-$144 km s-1 narrow absorption component of Na D. Based on STIS/HST UV
spectroscopy, Gull et al. (in prep.) noted similar disappearance of narrow
absorption lines. The HST/STIS spectra also show that the ionisation is
increasing not only in our LOS, but also within the SW Homunculus lobe,
changing dramatically to the NW of the stellar core. This indicates that the
coronagraph may be extended, with its edge projected against the LH. The
vanishing occulter modulating the FUV must be within the LH. The remaining
occulting material in the coronagraph also affects recombination lines, but at
larger distances, beyond the LH.

###  3.6 A circumstellar DIB at $\lambda$10792 Å?

Figure 6:  Profile of the absorption line at $\lambda$ 10792 Å. The black line
is from a CRIRES spectrum taken on 2009 April 03, showing at least four
components blueshifted and redshifted between $\pm$ 45 km s-1 from the central
reversal. Zero velocity corresponds to the position measured by Hamano et al.
(2015) for a diffuse interstellar band at the same position and the grey strip
marks 3 $\sigma$ uncertainty in their measurements. The magenta line shows an
observation at OPD/Coudé on 2021 March 08, with four times lower spectral
resolution. It seems that the two main components of the absorption line
changed the intensity ratio between 2009 and 2021.  Figure 7:  Diffuse
Interstellar Bands (DIBs) at $\lambda$10780 Å (DIB1) and $\lambda$10792 Å
(DIB2) in a sample of stars. Squares are from Hamano et al. (2015) and
triangles from this work. Points labelled in magenta indicate deviating
objects from the linear relation between the two DIBs: EW(10792 Å) $=$ (0.43
$\pm$ 0.03) EW(10780 Å). Blue triangles indicate a long-term fading of the
$\lambda$10792 Å absorption in $\eta$ Car (here called eta1.079).

We have explored the possibility that the isolated absorption feature at
$\lambda$ 10792 Å, hereafter eta1.079, is a circumstellar counterpart of a
diffuse interstellar band (DIB) (Díaz-Luis et al., 2015). In addition to the
agreement in wavelength position, its long-term strength decrease is well
correlated with the decreasing extinction, and it has a gradient opposite to
the P Cygni absorption strength of He i 10830Å and most of the resonant lines
in the optical range.

This feature was reported in $\eta$ Car by Damineli Neto et al. (1993) who
attributed it to a blue displaced absorption at RV $=$ $-$1040 km s-1
associated with the He i 10830Å line as also adopted by Groh et al. (2010).
However, no similar feature at this velocity was present in the He i 20587Å
“sister” line shown in that work and any other $\eta$ Car spectral line
reported so far.

In order to examine its profile and wavelength, we plot a high-resolution
spectrum reported by Groh et al. (2010) in Figure 6, taken on 2009 April 03.
Two other CRIRES spectra (from 2008 May 05 and 2009 Feb 09) reported in that
work show a similar line profile close to $\lambda$ 10792 Å except for a
central peak above the continuum, which might be a defect, not necessarily a
real emission. The absorption feature at this wavelength shows a complex
profile, with at least four components between $+$45 and $-$46 km s-1. There
is a central reversal feature in almost exact coincidence with the
$\lambda$10792 DIB reported by Hamano et al. (2015). The vertical gray strip
in Fig. 6 shows the 3 $\sigma$ uncertainty around $\lambda$10792.15 $\pm$ 0.15
Å reported by Hamano et al. (2015) for many directions in the Galaxy -
hereafter named DIB2.

A recent observation taken at the OPD/Coudé focus on 2021 March 08 with R $=$
15,000, four times lower than CRIRES, is also shown in Figure 6. The low
amplitude reversal is at the same velocity as in CRIRES. The two main
absorption components are shown at approximately the same velocity, but with
interchanged depth ratio, in addition to this variation, the EW has been
decreasing slowly over the long term, as described below.

Although the large positive velocity makes the identification of this feature
as DIB2 unlikely, in order to gain more information, we measured the DIB2 on
many normal and emission-line stars with spectra taken at the OPD/Coudé (Table
3). Figure 7 shows our measurements of DIB2 compared with the closest one at
$\lambda$10780 Å (hereafter DIB1). Figure 7 shows a very good linear
correlation for most of the observations of ours and of Hamano et al. (2015):
EW(DIB2)/EW(DIB1) $=$ 0.43 $\pm$ 0.03. These authors found a single
discrepancy of the correlation for CygOB2#12 and we show two additional
exceptions for the LBVs HD 316285 and W243 in the Westerlund 1 cluster. These
measurements are based on sufficiently large EWs, making them quite reliable.
These three objects have much larger extinction and larger DIB1 EW, but have
their DIB2 EW much fainter than expected. This indicates that DIB1 and DIB2
only mimic belonging to the same family of carrier molecules when measured in
the solar vicinity but might have different behaviour in other parts of the
Galaxy.

Figure 7 includes measurements of eta1.079 since 1995 (Table 4). The feature
eta1.079 is exceptionally strong compared to DIB1. When comparing its large EW
with CygOB#12, HD 316285 and W243, which deviate in the opposite sense with
respect to the linear fit, it seems unlikely that eta1.709 is produced by the
same carriers as DIB1. However, there are only a few cases studied at very
high reddening and we cannot reach a firm conclusion. Despite the relatively
large errors in EW, the long-term behaviour of this feature in $\eta$ Car is
clear: its intensity is decreasing with time. On the other hand, the strength
of DIB1 is consistent with constant EW. eta1.079 has been decreasing at a rate
of $\sim$ $-$10 mÅ yr-1 which evolves on the contrary sense to the P Cyg
absorption in many spectral lines. For example, the P Cyg absorption of He i
10830Å in the same period increased at a rate $\sim$ $+$ 500 mÅ yr-1. If
eta1.079 is not produced by DIB1 carrier molecules, it might be due to another
mixture of molecules and does not seem to be an absorption associated with the
primary wind.

###  3.7 Correlated decrease of $A_{\rm V}$ extinction and the eta1.079
absorption within the coronagraph

Figure 8:  Relation between eta1.079 ($\lambda$10792 Å absorption) and $A_{\rm
V}$. The extinction towards the stellar core ($x$-axis) was reported by
Damineli et al. (2019). It is fading in the form $A_{\rm V}$(mag) = 4.0 $-$
0.113 (year $-$ 2014.11) and is predicted to reach the ISM value $A_{\rm V}$
$\sim$ 1.5-2.4 mag in 2028-2036. The equivalent width of eta1.079 absorption
($y$-axis) is fading in tight correlation ($r$ = 0.997) with the visual
extinction of the central star: EW(DIB2) (mÅ) = 104.09 $A_{\rm V}$. The
typical error bar of the individual measurements is shown as a cross in the
top-left of the plot (see Table 4).

Damineli et al. (2019) showed that the rate of decrease in extinction follows
the equation $A_{\rm V}$ (mag) = 4.0 $-$ 0.113 (year $-$ 2014.11). If the
excess absorption of eta1.079 in our LOS to the central star originates in the
coronagraph, it would follow the extinction decrease. We used the above
relation to derive $A_{\rm V}$ for the dates we have measured EW(eta1.079) and
found a surprisingly tight correlation: EW(eta1.079) (mÅ) = 104.09 $A_{\rm V}$
(mag), with r $>$ 99% (dashed magenta line in Fig. 8). Spectra taken in the
time interval 1990-2021 are labelled for every decade in the plot to indicate
how the changes vary with time. Despite the relatively large errors in EW
measurements (see Table 4), we see that the changing EW of eta1.079 follows
the decreasing extinction inside the coronagraph.

The normal way to produce extinction is by dust, and so the correlation found
in Fig. 8 might suggest that the eta1.079 carrier molecules are associated
with dust grains or at least they are mixed inside the same region. The
brightening in the $V$-band between orbital cycles 10.74 (2002.07) and 13.69
(2018.44) was 1.86 mag as compared with 2.9 mag at 1200 Å measured in STIS
spectra (Damineli et al., in prep.). If the brightening was due to just
decreased extinction, the increase in brightening at 1200 Å should have been 8
magnitudes for a canonical reddening law ($R_{\rm V}$ = 3.1). This indicates
that the extinction of the occulter was mostly grey, which requires very large
dust grains or electron scattering in an ionised gas. Large dust grains are
supported by dust model results of Morris et al. (2017), in particular iron
grains in the range of few to tens of microns. Such large grains would
condense close to the central source at temperatures of a few hundred to 800 K
and are not in thermal equilibrium. Such a high temperature in principle is
against the identification of eta1.079 as the same carrier molecule as DIB2,
which is formed in the very cold ISM.

###  3.8 Spatial structures inside the coronagraph

Figure 9 compares H$\alpha$ and eta1.079 EWs plus $V$-band flux variability
from 1985 through 2020. The $y$-axis was scaled to represent the measured
absolute H$\alpha$ EW in Å. Both the eta1.079 EW and the $V$-band flux values
were multiplied by constants to fit within the plot. The fading of the
eta1.079 EW (black line) occurs in concert with the H$\alpha$ EW (red line).
The $V$-band flux of the whole object (stellar core + Homunculus, green line)
increases with time. A well-known change in the brightening rate occurred in
1997-99. It seems that there was a coeval fading of H$\alpha$ and eta1.079,
labelled as “transition” in Fig. 3.8. If this transition was as sharp as it
appears to be, it would not be consistent with systematic coronagraph
dissipation but implies a clumpy portion of the coronagraph moving out of our
LOS.

Figure 9:  Structures in the coronagraph. The long-term brightening of the
entire object (star + Homunculus) due to decreasing extinction towards the
stellar core is shown by the green line. Also shown are the EW of eta1.079
(black line) and the absolute EW of H$\alpha$ (magenta line), both measured in
the star. A strong inverse correlation between the EWs of eta1.079 and
H$\alpha$ with the $V$-band flux is observed. The $V$-band light curve (green
line) is based on data taken from Damineli et al. (2019) and transformed to
relative flux.

##  4 Discussion

###  4.1 A model for the central star

To facilitate the discussion we present a model for the primary star, assuming
a distance of 2.3 kpc to the object (Allen & Hillier, 1993; Meaburn et al.,
1993; Davidson et al., 2001; Shull et al., 2021).

Following previous works (Hillier et al., 2001; Hillier et al., 2006) we chose
as basic parameters of the model:
$R_{\hbox{\rm core}}\,=\,120$ $R_{\odot}$
$\dot{M}\,=\,8.0\times 10^{-4}$ $M_{\odot}\,\hbox{yr}^{-1}$
$L=5.0\times 10^{6}$ $L_{\odot}$
$V_{\infty}=420$ km s-1

These are similar, but not identical to earlier values adopted by Hillier et
al. (2001):
$R_{\hbox{\rm core}}=60$ $R_{\odot}$
$\dot{M}$ = 1.0 $\times$ 10-3 $M_{\odot}\,\hbox{yr}^{-1}$
$L$ = 4.0 $\times$ 106 $L_{\odot}$
$V_{\infty}$ = 400 km s-1

And also similar to Groh et al. (2012a):
$R_{\hbox{\rm core}}\,=60$ $R_{\odot}$
$\dot{M}$ = 8.0 $\times$ 10-4 $M_{\odot}\,\hbox{yr}^{-1}$
$L$ = 5.0 $\times$ 106 $L_{\odot}$
$V_{\infty}$ = 420 km s-1

The differences in the physical parameters are explained as follows. The
luminosity comes from infrared measurements - its precision (for a fixed
distance) is primarily limited by a correction for flux not absorbed by dust,
and a correction for the luminosity of the companion star. Due to the large
mass-loss rate, the primary wind is optically thick, even in the visible. The
core radius cannot be well determined - possible values range from 60 to
several hundred $R_{\odot}$. Larger values are preferred if we assume that the
He i line emissions arise primarily in the wind-collision interface between
the two stars as a consequence of the ionising flux of the secondary star. In
this work we adopt $R_{\hbox{\rm core}}\,=\,120$ $R_{\odot}$.

Models of the central star are riddled with systematic uncertainties arising
from the following:

  * •

Due to circumstellar extinction, we do not know the amount of reddening, nor
the shape of the local extinction law. As a consequence, there is no absolute
flux scale in the optical or UV.

  * •

As discussed below, the circumstellar extinction alters the spectral
appearance of the star (even with HST data). As a consequence, Hillier et al.
(2001), did not model the “intrinsic” spectrum of the central star in the
modelling of the HST spectrum obtained near periastron in December of 1998.

  * •

The companion star influences the observed spectrum via its ionising radiation
and the interaction of its wind with that of the primary (Groh et al., 2012a,
b; Madura & Groh, 2012). At apastron, the companion star lies at a distance of
$\sim$ 30 au ($\sim$ 6500 $R_{\odot}$) from the primary. The companion star
carves a cavity in the primary’s wind, and this cavity partially overlaps (in
radius) with the line formation region in the primary wind. Its influence on
the primary’s spectrum is partially governed by the opening angle of the
cavity, and hence the relative wind momenta of the two stars. While the
secondary is believed to be less luminous than the primary (e.g. Mehner et
al., 2010a), its ionising UV radiation field ($\lambda$ $<$ 912 Å) is much
stronger than that of the cooler primary.

  * •

Line emission may also arise from the wind collision zone.

  * •

The wind may be asymmetric (Smith et al., 2003b; van Boekel et al., 2003;
Weigelt et al., 2007).

Despite these difficulties and uncertainties, the model of the primary star
can still be used to glean insights into the nature of the coronagraph.

###  4.2 CMFGEN studies of impact on the spectrum by different obstructing
structures

Two models were constructed with the CMFGEN model (Hillier & Miller, 1998) to
assess the impact of an obstruction in front of the central star. First, we
computed the unobscured optical spectrum of $\eta$ Car A as shown by a blue
line in Fig. 10. The orange line in Fig. 10.a shows the spectrum integrated
outside a circular disk of radius $r$ $=$ 11.5 mas, concentric to the star. In
Fig. 10.b the magenta line shows the integrated flux from inside a slit of
infinite length and 25 mas wide. The slit is offset perpendicularly to its
length so that the centre of the star is 2.5 mas outside the slit. This second
geometry is inspired by the emission structure observed by Millour et al.
(2020) and by Falcke et al. (1996) and by the spatial variability of the
nebular ionisation analysed by Gull (in preparation). As expected, the
coronagraph enhances the wind lines relative to the unimpeded view of the
spectrum in both geometries, contributing to the majority of the continuum
flux. Since most of the Fe ii, and particularly the [Fe ii] emission lines are
formed at larger radii than the Balmer lines, they show a greater increase in
the EW absolute values. For the same reason, the intensity ratios of H$\alpha$
to the upper members of the Balmer series also increase.

Figure 10:  Spectra from CMFGEN models with and without obstruction of the
central star by a coronagraph. The blue line is the spectrum of the entire
object (star plus wind) Panel a: obstructing coronagraph as a circular disk
centered on the star with radius $r$ = 11.5 mas (26.45 au). Panel b:
extraction inside a slit of infinite length and 25 mas wide, in which the star
is placed at 2.5 mas (5.75 au) from the slit border. Insets display zoomed
views of spectral regions of interest.

In Fig. 10, excluding the Balmer lines, the strongest lines are from Fe ii,
while the weaker lines are due to [Fe ii]. The He i transition at $\lambda$
5876Å is much stronger in the full spectrum (blue curve) since the line forms
close to the star because Helium is mainly neutral in the outer wind.

There are numerous possibilities for the geometry and size of the coronagraph,
including sharp and smooth boundaries. Our intent was not to fit the observed
spectrum, which demands additional constraints from the observations, but to
explore the general EW trends of an occulter in front of the star and its wind
spectrum.

In Table 2 we present the EW for a set of lines calculated in three CMFGEN
model spectra shown in Fig. 10: the unobscured view (full); a coronagraph
occulting only the stellar core (disk); and a the occulter covering the whole
field of view, except for a slit, which are compared to two observed spectra:
one taken with FEROS/ESO in 1999.1, when the impact of the coronagraph was
still large ($A_{\rm V}\,\sim\,4$ mag), and one taken with NRES/LCOGT
spectrograph (with instrumental setup similar to that of FEROS/ESO) in 2019.9,
when the impact of the coronagraph had decreased to $A_{\rm V}\,\sim\,1$ mag.
The decrease in line strengths between these two dates follows the same trend
in the spectrum changing from that of an occulted half stellar wind (slit) to
the unocculted spectrum (full). This is clearly demonstrated in Table 2 which
shows EW as a function of the models (columns 2, 3 and 4) and for the
different epochs (columns 5 and 6).

Table 2: Equivalent Width (in Å) of representative lines extracted from CMFGEN
models and observations.

Feature | fullc | diskb | slita | 1999.1d | 2019.9d
---|---|---|---|---|---
H$\alpha$ |  $-$646 |  $-$3675 |  $-$2504 |  $-$876 |  $-$551
H$\beta$ |  $-$131 |  $-$355 |  $-$444 |  $-$139 |  $-$121
H$\delta$ |  $-$35 |  $-$81 |  $-$121 |  $-$28 |  $-$21
Fe ii 5317 |  $-$15 |  $-$178 |  $-$67 |  $-$13 |  $-$6

Notes: a extraction inside a slit aperture from $d$ = 2.5 to 27.5 mas (5.75 to
63.25 au) of the central star and infinite length; b extraction outside a
circle with $r$ = 11.5 mas (26.45 au); c spectrum extracted over the entire
field (star+wind); d observations taken with FEROS/ESO (1999.9) and NRES/LCOGT
(2019.9) instruments.

###  4.3 Observational constraints for the position and shape of the
coronagraph

While the dusty occulter may be changing by tangential motion, irregular,
clumped substructure and long-term dissipation, we describe a simple model to
constrain its size. Assume the coronagraph is circular, homogeneous,
motionless, concentric with the primary star, and only its transparency is
decreasing. The maximum size of the coronagraph can then be derived by using
the projected distance of the Weigelt knots (Weigelt & Ebersberger, 1986)
closest to the central star. While Weigelt B has faded, the two other clumps
(C and D) have remained constant in flux, at the same phase of the orbital
cycle, in the 1998-2004 period, in the optical and UV windows (Gull et al.,
2016). Unlike the central source, the Weigelt knots suffer very low
extinction, $A_{\rm V}$ $\sim$ 2.0 $\pm$ 0.5 mag (Teodoro et al., 2020). Since
the Weigelt C object was within $\approx$ 200 mas (462 au) from the central
source in 1995 (Weigelt & Kraus, 2012), we deduce that the radius of the
coronagraph must be smaller than this, assuming it to be spherical. The
coronagraph cannot be located between Weigelt C and the binary system since
the high-excitation spectrum of Weigelt C has remained constant over the last
two decades. The coronagraph seems not to have the same nature as the Weigelt
knot, because there are indication that it is more extended, covering the SE
direction. In addition, the Weigelt knots are located at Homunculus equator,
not to our LOS to the central binary.

An estimate of the minimum size along our LOS can be obtained from estimates
of the circumstellar extinction. In the modelling by Hillier et al. (2001),
the derived $V$-band circumstellar extinction was approximately 5 magnitudes
in 1998, which agrees with Damineli et al. (2019). With the physical model of
the coronograph described above, and assuming the coronagraph has infinite
optical depth, we obtain for a concentric circular configuration, a minimum
radius of 60 au (26 mas). This is roughly the radius at which the electron
scattering optical depth is 0.01. Since the spectrum extracted beyond this
radius shows distinct differences from that of the primary star, the
coronagraph must have some transparency, and thus it must have been larger
than 60 au in 1998. However, its size is not well constrained to date. The
flux increase observed since 1998 could be due to either an increase in
transparency, a decrease in radius or a shift in the coronagraph location
relative to our LOS.

Of course, the coronagraph does not have to be circular. Falcke et al. (1996)
performed speckle polarimetry and resolved an obscuring $\sim$ 20 $\times$ 60
mas SW-NE dust bar in polarised light. Recently, Millour et al. (2020)
reported adaptive optics observations of $\eta$ Car obtained with SPHERE at
the VLT/ESO. The spatial resolution was 20 mas ($\sim$ 47 au), and
observations were obtained in both H$\alpha$ and the nearby continuum. While
difficult to interpret unambiguously, the images show structures that might be
connected to the coronagraph. The stellar core is occulted by a dust lane
running SW to NE and there is a bi-conical structure in the polarisation map
running perpendicular to it (their Fig. 4). Further, the area NW of the
central star is more transparent than the opposite side, consistent with
markedly less extinction to Weigelt C and D.

###  4.4 Consequences of the disappearance of the coronagraph

The absolute EW of the eta1.079 feature is fading at a rate of $-$29 mÅ yr-1,
so it will likely disappear in $\sim$ 2030\. The uncertainty in this
prediction is large because this feature is faint and the spectra are noisy,
but it is consistent with the prediction for the coronagraph vanishing time-
scale derived from the decreasing rate of the $A_{\rm V}$ extinction.
Following Damineli et al. (2019) (their Fig. 18) this will happen in the time
frame 2028-2036, depending on the foreground extinction, which is uncertain in
the range $A_{\rm V}$ = 2.4-1.5 mag.

The disappearance of the coronagraph is expected to have an impact also on the
absorption lines of the spectrum components, since it will allow the light
from the secondary star to ionise far-out regions of the circumstellar medium
that were previously shielded from it as we have shown for the narrow
H$\alpha$ absorption in the LH. The details will be evaluated only after the
size and position of the coronagraph are known.

The correlation shown in Fig. 8 indicates that the eta1.079 carrier molecules
are related to dust grains and both are being photo-dissociated by the intense
stellar radiation field. When the coronagraph has completely vanished, the
stellar core will stop increasing in brightness. The predicted magnitude of
the stellar core at that time is $V$ $\sim$ 2.5-3.5 mag, which is a little
brighter than it was in the 1600sbf (the amount of light reflected from the
Homunculus is not very relevant here, because it was absent before the GE and
now is fainter than the central star). The true luminosity variation depends
on how much the bolometric correction (BC) changed since then, but
unfortunately there were no historical reports on the colour of the star
before the GE event. Let us assume that the change in BC was modest, which
implies that the intrinsic luminosity did not change significantly or might
have mildly increased. Since the GE ejected $>$ 15 M⊙ following Morris et al.
(1999) and Smith et al. (2003a) or $>$ 45 M⊙ following Morris et al. (2017) \-
using a full radiative transfer model with a far more chemically realistic
dust library - the luminosity of the central star should have dropped by a
large amount, unless it was compensated by some process. This favours the
scenario of a merger in a triple system (Smith et al., 2018; Portegies Zwart &
van den Heuvel, 2016; Hirai et al., 2020) over a massive wind ejection as the
mechanism for the GE, since the primary star would have been mixed internally
in a merger, bringing fresh fuel to its central core.

##  5 Conclusions

This work brings together various pieces of evidence indicating the presence
of an obscuring structure in front of the binary system in $\eta$ Car. Some of
them were suspected as long as 30 years ago but now there are better
observational data and modelling.

The central stellar binary ($\eta$ Car A + $\eta$ Car B) has been brightening
by many magnitudes for half a century, without major changes in the intrinsic
spectrum of the primary star. This is not expected for a star which was
already close to the Eddington limit at the beginning of this process. The
long-term brightening has a circumstellar origin. The Homunculus reflection
nebula has remained at an essentially constant brightness ($V$ $\sim$ 5.48
mag) while the central stellar core has brightened by several apparent
magnitudes (Damineli et al., 2019).

The main conclusions of this spectroscopic study are:

  1. [i]

  2. 1.

Spectral lines formed in the stellar wind have larger absolute EW in direct
than in reflected light (Fig. 1).

  3. 2.

The absolute EW values of the lines formed in the wind have decreased
systematically from cycle to cycle (Fig. 2.a). However, the Weigelt knots
located at a projected distance of a fraction of an arcsec from the central
star have maintained a constant brightness and ionisation level (Fig. 2.b).

  4. 3.

After correcting for the secular brightening, the broad lines formed in the
stellar wind exhibit a constant line profile (Fig. 2.c).

  5. 4.

Line profiles of the upper Balmer series members have been almost constant for
the last 20 years, as compared to large variations in H$\alpha$, Fe ii and [Fe
ii]. This is explained by the fact that these lines are formed at large radii
in the stellar wind, as compared with the high-energy spectral lines, which
are mainly formed behind the coronagraph (Fig. 3). The absolute EW value of
H$\alpha$ in direct light has decreased by $\approx$ 50% from 1985 to 2014
(Fig. 4). In contrast, the EW of H$\alpha$ observed in the FOS4 reflected
light remained roughly constant over the time interval 2000-2010. The
H$\alpha$ EW in direct light converged to a value close to that from the FOS4
reflected light in 2014 and stalled afterwards (Fig. 4). EWs from both views
remained at the same level during the last orbital cycle, suggesting that the
coronagraph might have stopped vanishing. Alternatively, the coronagraph may
have fully vanished, and this has happened $\sim$ 10 years earlier than
previously predicted (Damineli et al., 2019). Spectra produced by CMFGEN
modelling in the presence or not of an artificial coronagraph of varying
sizes, shapes and positions, are in general agreement with observations (Fig.
10 and Table 2)

  6. 5.

The constancy of the H$\delta$ and H$\alpha$ EWs at the FOS4 reflected light
reveals that the primary star in $\eta$ Car is more stable than has often been
claimed, for example by Mehner et al. (2012). This indicates that whatever
occurred during the GE, after $\sim$ two centuries the star has become stable
again - see (Figs. 3 and 4).

  7. 6.

The variability of the coronagraph is not completely smooth as seen in the
short-lived event that occurred at the end of the 1990’s. This could be due to
irregularities in the transparency of a moving occulter, or a temporary
variation in the rate of vanishing (Fig. 9).

  8. 7.

The variability curve of the H$\alpha$ narrow absorption (“shell”) line (Fig.
5) has a complex pattern, indicating that it is produced by a number of
structures that may be moving across our LOS. All of these structures produce
an ionisation shadow in the LH at the position crossed by our LOS to the
binary system. The non-periodic component of H$\alpha$ narrow absorption
(“shell”) in Fig. 5 is produced by the coronagraph and it is weakening over
the long term, in line with other evidence.

  9. 8.

In addition to the coronagraph, there is a second kind of occulter that
crosses our LOS to the binary system around periastron, when the ionising
secondary star is on the back side. This occulter is the primary’s wind, which
casts an ionisation shadow to the LH at our LOS close to periastron. It
produces the periodic component displayed in the variability curve of the
H$\alpha$ narrow absorption.

  10. 9.

The fact that the disappearance of the coronagraph impacts the strength of
absorption lines originating in the LH indicates that it is located inside the
LH.

  11. 10.

Our best guess is that the coronagraph is a sheet-like a circumbinary disk
seen at some inclination to our LOS - the NW border of which runs over the
central stellar object and is extended to the SE as suggested by the Millour
et al. (2020) and by Falcke et al. (1996) high spatial resolution polarimetric
maps. The permanent H$\alpha$ shell absorption reported by Boumis et al.
(1998) could be produced by the same disk in the Homunculus equator. Partial
occultation to the SE direction is suggested by the slow EW fading of Fe
ii$\lambda$4585 reflected at FOS4. This evidence points to a 3D geometry of
the occulter.

  12. 11.

We could not identify the origin of the absorption feature at $\lambda$10792
Å. It is unlikely to be formed in the primary’s stellar wind, because there is
no corresponding feature in any other spectral line of $\eta$ Car. It does not
seem to be produced by the same DIB known in the ISM, since it has positively
displaced absorption components to our LOS ($+$ 45 km s-1) and ISM DIBs are
produced at much colder temperatures. Because its decreasing EW is tightly
correlated with the decreasing extinction, we believe it is located inside the
coronagraph, produced by an unidentified process/molecule.

The different molecular absorptions, with a range of velocities and
equilibrium temperatures observed by ALMA indicate that there is a
contribution of absorbers from a range of distances from the central star in
our direction. Regarding the end of the changes in the coronagraph, if it will
continue up to the year 2030, the brightness of the central star will surpass
that of the situation previous to the GE. If this happens, taking into account
the large quantity of mass shed during the GE, the primary star should have
been rejuvenated and increased its luminosity. This would be in favour of the
scenario of a binary merging in a triple stellar system. However, the question
of how much such an internal mixing process could compensate for the lost mass
can only be tackled by a specific model.

Significant progress to map the coronagraph could be made with high spatial
resolution ($\sim$ mas) spectroscopy by sampling molecular absorption lines or
resonant emission-lines.

## Acknowledgements

AD thanks to FAPESP for support award numbers 2011/51680-6 and 2019/02029-2.
FN acknowledges FAPESP (2017/18191-8). AFJM is grateful for financial aid to
NSERC (Canada). MFC was supported by NASA under award number 80GSFC21M0002. We
thank the referee A. Mehner for many constructive criticisms.

## Data availability

The data underlying this article will be shared on reasonable request to the
corresponding author.

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##  Appendix A Extended data tables

Table 3: DIB1 and DIB2 Equivalent widths in normal and emission-line stars Object | EW $\lambda$10780 | $\sigma_{\rm{EW}}$ | EW $\lambda$10792 | $\sigma_{\rm{EW}}$ | Spectral Type
---|---|---|---|---|---
| (mÅ) | (mÅ) | (mÅ) | (mÅ) |
$\beta$ Monocerotis | 72 | 10 | 66 | 10 | B4Ve
$\delta$ Scorpii | 180 | 20 | 136 | 18 | B0.3IVe
$\beta$ Lyrae | 82 | 12 | 56 | 9 | B8.5Ib-II
HD 169454 | 169 | 18 | 86 | 11 | A1V
AG Carinae | 130 | 15 | 66 | 9 | WN11h-LBV
$\chi$ Ophiuci | 94 | 13 | 46 | 8 | B2Vne
HD 168625 | 242 | 30 | 110 | 15 | B6Iap-LBV
HR Carinae | 196 | 27 | 89 | 16 | LBV
HD 168607 | 221 | 18 | 74 | 10 | B9Iaep-LBV
MWC 314 | 267 | 30 | 87 | 12 | B3Ibe-LBV
P Cygni | 91 | 16 | 29 | 7 | LBV
HD 152236 | 61 | 16 | 19 | 7 | B1Ia-0ek
HD 316285 | 387 | 100 | 90 | 20 | B0Ieq-LBV
WD1-W243 | 690 | 120 | 134 | 30 | LBV

Table 4: Equivalent widths of the $\lambda$10792 Å absorption in eta Carinae observed at OPD/Coudé Year | EW | $\sigma_{\rm{EW}}$ | Year | EW | $\sigma_{\rm{EW}}$ | Year | EW | $\sigma_{\rm{EW}}$
---|---|---|---|---|---|---|---|---
| (mÅ) | (mÅ) |  | (mÅ) | (mÅ) |  | (mÅ) | (mÅ)
1990.068 | 650 | 66 | 2001.446 | 550 | 52 | 2010.177 | 490 | 47
1991.398 | 650 | 100 | 2002.325 | 520 | 93 | 2010.498 | 500 | 109
1993.32 | 630 | 68 | 2002.326 | 530 | 57 | 2011.196 | 490 | 39
1993.553 | 630 | 40 | 2002.331 | 520 | 45 | 2011.53 | 440 | 45
1995.36 | 660 | 68 | 2002.547 | 540 | 35 | 2011.531 | 500 | 52
1996.075 | 660 | 45 | 2003.360 | 540 | 52 | 2012.504 | 420 | 40
1996.346 | 640 | 66 | 2004.502 | 520 | 103 | 2013.537 | 450 | 44
1997.381 | 630 | 75 | 2005.567 | 500 | 48 | 2014.582 | 401 | 39
1997.482 | 600 | 82 | 2006.199 | 530 | 55 | 2014.585 | 398 | 57
1998.054 | 610 | 58 | 2006.429 | 500 | 42 | 2014.591 | 384 | 41
1998.103 | 620 | 49 | 2006.637 | 530 | 99 | 2014.593 | 400 | 36
1998.284 | 620 | 90 | 2007.491 | 520 | 54 | 2014.596 | 407 | 48
1999.125 | 620 | 59 | 2007.494 | 530 | 85 | 2016.486 | 387 | 75
1999.316 | 540 | 84 | 2008.381 | 500 | 52 | 2017.595 | 402 | 49
2000.132 | 530 | 50 | 2009.131 | 510 | 88 | 2018.417 | 380 | 42
2000.945 | 530 | 44 | 2009.134 | 550 | 101 | 2019.322 | 340 | 34
2001.435 | 550 | 35 | 2009.301 | 480 | 38 | 2019.542 | 360 | 32
2001.438 | 540 | 57 | 2009.430 | 510 | 110 | 2021.218 | 334 | 37

Table 5: Equivalent widths in direct light for H$\alpha$ (first 10 rows). The full table is available online as supplementary material and in the CDS. HJD | EW H$\alpha$ | $\sigma_{\rm{EW}}$ | Site
---|---|---|---
| (Å) | (Å) |
49063.609 |  $-$1009 | 48 | ESO/FEROS
49065.609 |  $-$1013 | 51 | ESO/FEROS
49066.621 |  $-$996 | 49 | ESO/FEROS
49067.605 |  $-$1005 | 50 | ESO/FEROS
49068.609 |  $-$1001 | 52 | ESO/FEROS
49069.609 |  $-$1009 | 50 | ESO/FEROS
49071.613 |  $-$1006 | 50 | ESO/FEROS
49072.613 |  $-$962 | 52 | ESO/FEROS
49073.648 |  $-$1004 | 49 | ESO/FEROS
49074.625 |  $-$998 | 49 | ESO/FEROS
Table 6: Equivalent widths in direct light for Fe ii$\lambda$4585 (first 10 rows). The full table is available online as supplementary material and in the CDS. HJD | EW Fe ii$\lambda$4585 | $\sigma_{\rm{EW}}$ | Site
---|---|---|---
| (Å) | (Å) |
48844.4880 |  $-$16.1 | 1.8 | ESO/FEROS
49063.6090 |  $-$15.2 | 1.7 | ESO/FEROS
49067.6050 |  $-$16.0 | 1.7 | ESO/FEROS
49069.6090 |  $-$16.0 | 2.8 | ESO/FEROS
49071.6130 |  $-$16.1 | 1.5 | ESO/FEROS
49072.6130 |  $-$15.0 | 3.2 | ESO/FEROS
49073.6480 |  $-$15.3 | 1.3 | ESO/FEROS
49075.5740 |  $-$15.2 | 2.3 | ESO/FEROS
49076.6250 |  $-$15.0 | 2.7 | ESO/FEROS
49077.5900 |  $-$15.4 | 2.3 | ESO/FEROS
Table 7: Equivalent widths in direct star for H$\delta$ (first 10 rows). The full table is available online as supplementary material and in the CDS. HJD | EW H$\delta$ | $\sigma_{\rm{EW}}$ | Site
---|---|---|---
| (Å) | (Å) |
58502.451 |  $-$27.4 | 1.0 | LCO/NRES
59039.234 |  $-$25.4 | 2.0 | LCO/NRES
59051.232 |  $-$26.4 | 0.6 | LCO/NRES
59152.602 |  $-$24.1 | 1.4 | LCO/NRES
59155.600 |  $-$25.0 | 1.2 | LCO/NRES
59174.562 |  $-$24.6 | 0.8 | LCO/NRES
59206.513 |  $-$23.8 | 1.1 | LCO/NRES
59210.463 |  $-$24.5 | 0.9 | LCO/NRES
49063.609 |  $-$28.8 | 13.5 | ESO/FEROS
49065.609 |  $-$32.5 | 12.2 | ESO/FEROS
Table 8: Equivalent widths of H$\alpha$ H$\delta$ and Fe ii$\lambda$4585 from UVES observations extracted at FOS4 position. JD-2400000 | year | EW H$\delta$ | $\sigma_{\rm{EW}}$ | JD-2400000 | year | EW Fe ii$\lambda$4585 | $\sigma_{\rm{EW}}$ | JD-2400000 | year | EW H$\alpha$ | $\sigma_{\rm{EW}}$
---|---|---|---|---|---|---|---|---|---|---|---
|  | (Å) | (Å) |  |  | (Å) | (Å) |  |  | (Å) | (Å)
51533.832 | 1999.97 |  $-$22.00 | 0.61 | 51533.832 | 1999.97 |  $-$7.00 | 0.49 | 52616.833 | 2002.94 |  $-$521 | 8
52634.854 | 2002.99 |  $-$20.00 | 0.62 | 52634.854 | 2002.99 |  $-$5.64 | 0.41 | 52639.852 | 2003.00 |  $-$560 | 9
52639.840 | 2003.00 |  $-$21.00 | 0.57 | 52662.862 | 2003.06 |  $-$6.92 | 1.38 | 52642.832 | 2003.01 |  $-$555 | 9
52662.862 | 2003.06 |  $-$19.60 | 0.61 | 52674.862 | 2003.09 |  $-$6.90 | 0.21 | 52658.814 | 2003.05 |  $-$530 | 9
52662.862 | 2003.06 |  $-$19.53 | 0.61 | 52990.816 | 2003.96 |  $-$5.90 | 0.36 | 52662.848 | 2003.06 |  $-$532 | 9
52968.807 | 2003.90 |  $-$20.04 | 0.73 | 53029.774 | 2004.07 |  $-$5.55 | 0.38 | 52674.848 | 2003.09 |  $-$530 | 9
52990.816 | 2003.96 |  $-$19.93 | 0.70 | 53055.653 | 2004.14 |  $-$6.54 | 0.43 | 52684.657 | 2003.12 |  $-$518 | 9
53029.774 | 2004.07 |  $-$20.33 | 0.67 | 53075.562 | 2004.19 |  $-$6.91 | 0.46 | 52990.803 | 2003.96 |  $-$564 | 6
53055.653 | 2004.14 |  $-$21.32 | 0.67 | 53389.738 | 2005.05 |  $-$6.50 | 0.58 | 53029.763 | 2004.07 |  $-$540 | 7
53075.562 | 2004.19 |  $-$20.56 | 0.93 | 53431.818 | 2005.17 |  $-$6.18 | 0.48 | 53075.703 | 2004.19 |  $-$564 | 9
53389.738 | 2005.05 |  $-$19.08 | 0.82 | 53866.524 | 2006.36 |  $-$6.68 | 0.51 | 53349.850 | 2004.94 |  $-$535 | 11
53431.818 | 2005.17 |  $-$19.50 | 0.66 | 53912.550 | 2006.48 |  $-$6.84 | 0.55 | 53389.709 | 2005.05 |  $-$540 | 14
53866.524 | 2006.36 |  $-$19.88 | 0.72 | 54535.755 | 2008.19 |  $-$6.67 | 0.41 | 53389.725 | 2005.05 |  $-$541 | 8
53912.550 | 2006.48 |  $-$21.18 | 0.76 | 54554.759 | 2008.24 |  $-$5.62 | 0.44 | 53431.805 | 2005.17 |  $-$527 | 26
54513.829 | 2008.13 |  $-$21.32 | 1.24 | 54567.535 | 2008.28 |  $-$5.20 | 0.66 | 53912.538 | 2006.48 |  $-$540 | 13
54535.755 | 2008.19 |  $-$20.10 | 0.58 | 54583.543 | 2008.32 |  $-$5.70 | 0.41 | 54475.731 | 2008.03 |  $-$540 | 12
54554.759 | 2008.24 |  $-$21.00 | 0.67 | 54599.492 | 2008.36 |  $-$5.24 | 0.44 | 54554.740 | 2008.24 |  $-$523 | 13
54567.535 | 2008.28 |  $-$21.79 | 0.70 | 54629.491 | 2008.45 |  $-$5.27 | 0.38 | 54567.517 | 2008.28 |  $-$559 | 13
54583.543 | 2008.32 |  $-$20.88 | 0.64 | 54657.552 | 2008.52 |  $-$5.35 | 0.47 | 54583.529 | 2008.32 |  $-$543 | 14
54599.492 | 2008.36 |  $-$20.46 | 0.63 | 54999.565 | 2009.46 |  $-$4.34 | 0.39 | 54599.475 | 2008.36 |  $-$552 | 14
54629.491 | 2008.45 |  $-$20.08 | 0.67 | 55013.483 | 2009.50 |  $-$4.86 | 0.34 | 54616.536 | 2008.41 |  $-$547 | 13
54657.552 | 2008.52 |  $-$20.20 | 0.66 | 56620.846 | 2013.90 |  $-$4.00 | 0.71 | 54629.504 | 2008.45 |  $-$529 | 13
54999.565 | 2009.46 |  $-$20.05 | 0.67 | 56650.835 | 2013.98 |  $-$4.31 | 0.71 | 56999.838 | 2014.94 |  $-$485 | 36
55013.483 | 2009.50 |  $-$20.00 | 0.44 | 56684.722 | 2014.07 |  $-$4.05 | 0.61 | 56999.841 | 2014.94 |  $-$487 | 12
56620.846 | 2013.90 |  $-$21.00 | 1.26 | 56712.686 | 2014.15 |  $-$4.28 | 0.74 | 57012.805 | 2014.97 |  $-$505 | 13
57043.865 | 2015.06 |  $-$19.08 | 1.39 | 57043.865 | 2015.06 |  $-$3.68 | 0.82 | 57027.837 | 2015.01 |  $-$497 | 27
57047.796 | 2015.07 |  $-$19.63 | 0.99 | 57047.796 | 2015.07 |  $-$3.87 | 0.59 | 57047.781 | 2015.07 |  $-$531 | 13