SAG: Stars and Galaxies Vol. 6, id. 1, 2023 December 28
c 2023. Center for Astronomy, University of Hyogo.
⃝

Long-term evolution of the Hα emission line of Pleione
between Jan 2009 and Mar 2023 in the BeSS database
Junichi Katahira1∗
1

Sakai-city, Osaka 591-8021, Japan∗
atak@db3.so-net.ne.jp(Received 2023 September 29; accepted 2023 November 10)

Abstract
Using the Hα double peak emission lines of Pleione in the BeSS database from January 2009 to March 2023, wehave measured the characteristic three point intensities of the line: the intensities at the blue and red emission peaks,
and the intensity of the central minimum between the peaks, and calculated the radial velocities of these three points.
These values are shown over a period of 14 years. The long-term variations shown are interpreted to be split into (a)the short-term variations when the companion of Pleione is close to periastron, and (b) the long-term variations in
the background. For the (b) variation, we try to explain the trajectory of the central minimum intensity by the disk
precession, and discuss the occurrence of the V/R phenomenon.
Key words: stars: Be star — stars: Pleione — stars: Be disk

1. Introduction
Pleione (28 Tau, BU Tau, HD23862) is a bright member of
the Pleiades cluster, with a magnitude of V = 5.09 (SIMBAD)
and a spectral type of B8Vne (Taranova et al. 2008), and started
in December 2005 the third cycle of disk formation activity
with a period of about 35 years (Tanaka et al. 2007), whichcurrently lasts until October 2023. Each disk-forming activity begins with a rapid decrease in stellar magnitude ∆V ∼ 0.5
mag, and exhibits spectral phase changes from the Be phaseto the Be-shell phase. During a disk-formation activity, the respective periods for the Be-shell and the Be phases are curiosly
separated by 17 years apart (Hirata 1995). Pleione has a double
peaked emission proﬁle of Hα, and during the Be-shell phase,shows a deep shell absorption with a minimum below the continuous light level. Only a few Be stars show such a deep shellabsorption. Pleione is also characterised by a periodic variation in the intensity ratio V/R of the blue and red emission
peaks during the disk formation activity period (Okazaki 1997;
Hirata 2007).
Pleione was detected as a speckle binary (McAlister et al.
1989), and Gies et al. (1990) conﬁrmed its companion during the lunar occultation of Pleione. Subsequent observations
using adaptive optics have detected more distant companions(Roberts et al. 2007). A hypothesis has been proposed that the
speckle companion approaches the host star on a 34-year cycle,
causing a disk-forming activity period (Gies et al. 1990), but it
appears that the analysis of the orbital elements of the speckle
companion has not progressed easily. On the other hand, theanalysis of the orbital motion at the radial velocity indicatedby the disk absorption line and the Hα emission line wing preceded this speckle analysis, and it was clariﬁed that Pleione isa single-lined spectroscopic binary star with an orbital motion

period of 218 days and an eccentricity e = 0.6 – 0.7 (Katahira
et al. 1996; Nemravova et al. 2010). Thus, Pleione is currently
considered to be at least a quadruple binary.
Since the early 2000s, a group of European amateur spectroscopic observers has published the middle-dispersion spectraof Pleione with high temporal resolution in the BeSS database1 .
Pollmann and Vollmann (2014) showed that each of the three
periastron passes of the companion is accompanied by drastic variations in the Hα proﬁle (such as large ﬂuctuations in
the V/R). These observations have visualised the interaction
between the disk and the companion in a single-lined spectroscopic binary system with high eccentricity (hereafter referred
to as the periastron phenomenon).
This article reports on the re-measurement of the attractive
observational data of Pleione in the BeSS database. In section
2, we describe the data selection from the BeSS database. In
section 3, we explain the data processing and the measurement
method of the Hα proﬁle, and show the long-term variations forthe intensities at the blue and red emission peaks (I B, I R)
and at the central minimum (I c) between the peaks, and for theradial velocities calculated at the above three points. We separate the long-term variation into the short-term variation andthe long-term background variation in order to discuss the evolution of the disk. In section 4, we try to interpret the long-term
background variation of I c by the disk precession of Hirata(2007). In the V/R long-term background variation, we show
the possibility that the ”V/R phenomenon” is triggered. The
coexistence of short-term V/R variations in the periastron phenomenon and the ”V/R phenomenon” in the long-term background variation is also touched upon.

1

BeSS database ⟨http://basebe.obspm.fr/basebe/⟩

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J. Katahira

2. Data Selection
The data in the BeSS database consist of observations madewith a variety of instruments (aperture diameter:10 – 40 cm,
long-slit spectrographs, echelle spectrographs, and various
CCD detectors). These data have been already analysed by
the group of observers introduced in section 1. The following
points have been taken into account in order to deal with thelong-term data again. For the study of long-term variations, it
is desirable that the wavelength resolution R is constant, sincemeasurements of the central intensity of the emission line proﬁle, etc., can be obtained as a function of R. On the other hand,to increase the time resolution of the observation, it is necessary to select observations with a considerable range of R. To
satisfy both these requirements, observations with R ranging
from 10000 to 20000 were selected. In addition, it is requiredthat the trend of changing proﬁle intensity indicated by individual observations is relatively stable. As a result, a total of
520 observations were obtained. The average R-value distribution of the selected observations is mostly around 15000. Notethat the R value follows the value described in the header of thedownloaded ﬁts ﬁle.
3. Analysis and Results
3.1. Measurement work
Each observed spectrum was normalised to the continuum
light level (using Nijiboshi software2 ) and converted to the heliocentric wavelength system (use of IRAF). The processed
spectra of individual observations can be downloaded from thepersonal website3 .
The parameters for the intensity and radial velocity at threecharacteristic points of the normalised Hα proﬁle, i.e. the intensities I B, I R, radial velocities V r B, V r R at the blueward and redward emission peaks, and the intensity I c, radial velocity V r c at the central minimum, are calculated by
the quadratic least squares ﬁt to the above three points of theline. Note that the proﬁle near the emission peak shows very
complex variations, so we work with a small number of data
points (about 10 – 20 data points) around the measurement
point. Also, if the H2 O absorption components of the Earth’s
atmosphere are poorly mixed, the emission peak data are notmeasured (only the I c and V r c values are measured). The
continuum of the Hα emission line is estimated by interpolating from the continuum light level of the spectrum in the
wavelength range widely sandwiched around the Hα line (interpolation with a cubic spline function). Since the adjustment
error of the continuum light is about ± 1%, the calculation
error of the three intensity parameters of the emission line isabout the same, reﬂecting the former, but the observation errorof each observation data is quite large. Assuming that the observation date outside the period of periastron phenomena has
a constant value, the empirical observation error at the emission peaks where atmospheric absorption tends to overlap to be
δI B and δ I R = ± 0.07, and at the central minimum intensity
is estimated to be δ I c = ± 0.05 (the magnitude of the error in
2
3

Nijiboshi software ⟨http://yamanecoland.starfree.jp/nijibosi.html⟩
personal website ⟨https://bessdata.yu-yake.com⟩

Vol. 6,ﬁgure 1). Measurements separated by more than this error are
considered reliable. Similarly, the empirical errors in radial velocity observations are δ V r B and δ V r R = ± 5 km s−1 andδ V r c = ± 5 km s−1 (the magnitude of the error in ﬁgure 2).
In addition, the errors in the emission line intensity ratio
V/R (= I B / I R) and the mean emission peak radial velocity
V r EP = 0.5 × (V r B + V r R) derived from these measurements are empirically ± 0.02, ± 3 km s−1 (the magnitude ofthe error in ﬁgure 3).
3.2. Intensity parameters (I B, I R, I c)
Figure 1 (top) shows the variation of the three intensity parameters for the blue and red emission peaks (I B, I R) andthe central minimum (I c) from January 2009 to March 2023.
The horizontal time scale is the heliocentric median time ofthe observation (header information of the spectrum) minus
245000. This horizontal time scale will be referred to simply
as ”T ” in the following. The green vertical dashed line in theﬁgure indicates when the companion becomes periastron, andthe time is calculated according to the orbital parameter solution 3 of Nemravova et al. (2010). This is also the time when
the radial velocity of the primary star due to orbital motion isalmost minimal.
Figure 1 shows that the state of the variation discussed by
Pollmann (2018, 2020) is reproduced. Also of interest is the
rapid rise in I c above continuous light levels and the rapid increase in emission line peaks at the periastron phenomenon.
During the period of T = 9800 – 9900, the observations forwhich the three data are not complete have the damage at theemission peaks due to the strong atmospheric H2 O absorption. Their emission peaks (I B, I R) were not measured.
However, during the same period, there is a characteristic variation in the I c value, so they are specially displayed.
In January 2009, at the left end of the ﬁgure, according to
Tanaka et al. (2007), the Hα emission line component originating from the old disk disappeared, and only the emission fromthe new disk was observed. Thus, it can be said that the ﬁgure is an observational record of the evolution of the new disk.
The ﬁgure appears to superimpose the short-term variations on
top of the long-term background variation for the three parameters. It is clear that the onset of each short-term variation coincides approximately with the periastron time. Pollmann (2020)shows almost the same value as the orbital period of Pleione by
Nemravova et al. (2010) from the periodic analysis of the I cvalue variations. Thus, repeated short-term variations are variations caused by the approach of the companion star to the diskof the primary star (periastron phenomenon). The long duration of the short-term variations suggests that the gravitational
pull of the companion is inducing structural density changes inthe disk.
3.2.1. Long-term background variation in I c
If we can separate the short-term variations, which periodically overlap in ﬁgure 1, from the long-term variations, it will
be easier to understand the parameter changes in the ﬁgure.
We ﬁrst describe how to determine the long-term backgroundvariation of the I c value, which appears to be a simple change.
Based on the distribution map of the I c value seen in the phaseof the binary orbital period, and following the tracing of theHα proﬁle time series, the periastron phenomenon of the short-Long-term evolution of the Hα line of Pleione

Intensity

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Fig. 1. Long-term intensity variation of the blue peak (I B), red peak (I R), and central minimum (I c) of the Hα double peak emission line. The
horizontal scale subtracts 245000 from the heliocentric median time JD of observation. The green vertical dashed line indicates the periastron time of
the companion star. In the sequence of ”I c”, data outside the periastron phenomenon are marked with blue circles. These data are ﬁtted by the leastsquares method, and the result is plotted by the dashed lines, which is considered to be the long-term background variation. See the text for details. (Top)All data from January 2009 to March 2023. (Bottom) Enlargement after July 2018. After T = 9000 the pattern of the long-term background variation is
tentatively highlighted by the dashed purple line.

4

J. Katahiraterm variation continues from about 15 days before periastronto 100 days after the periastron time (in the sequence of ”I c”
in ﬁgure 1, data outside this period are marked with blue circles). In fact, this period is more than half of the binary orbital
period of 218 days, but the stable I c value continues outside
this period (but, the range used is up to T = 9000). From thisobservational estimate, it can be assumed that the long-term
background variation appears in the observed quantities during the period when the short-term variations disappear. The
I c values marked with blue circles show a trace of the longterm background variation. These points are ﬁtted by the leastsquares method, and the results are shown in the ﬁgure by theblue dashed line (for the quadratic formula, correlation coefﬁcient 0.90) and the red dashed line (for the cubic formula,correlation coefﬁcient 0.91). The changes shown by the two
curves are considered to be the long-term background variations of the I c value between T = 4800 and T = 9000. The
curves are extended to T = 9500 to compare the variation after
T = 9000. The above approximation method for estimating thelong-term background variation is only valid up to T = 9000
because it combines a small (visible to be stable) part of thelong-term variation.
From the long-term background variation of I c obtained,
let’s ﬁrst consider the characteristics of the changes in the disk.
It also provides a clue to consider the long-term variation of
the emission line peak intensity in the ﬁgure. The resulting
I c long-term background variation curve (blue dashed line)
starts to decrease from around I c = 0.72, down to the minimum I c = 0.21, and then starts to increase again, reachingI c = 0.63. This variation gives the impression of almost symmetrical changes centred on the minimum. As I c < 1 in thelong-term background variation, the spectral phase due to theshape of the Hα emission line is the Be-shell phase. The timeof minimum in the I c curve is estimated to be around T = 7000
(2014 – 2015). This epoch corresponds to the maximum shell
activity in Hα.
Looking at the evolution of the emission peak intensitychange in the period outside the periastron phenomenon (whenthe long-term background variation is visible) in ﬁgure 1, theemission intensity shows an upward trend. However, it should
be noted that after the maximum of the shell activity at about
T = 7000, the increase of the emission intensity stagnates for awhile. This tendency may be due to the effect of the maximum
shell activity.
The Fujii Bisei Observatory4 has published a long-term variation diagram of the equivalent width of the Hβ line and of themetallic absorption equivalent width regionally summed over600 Å in Pleione5 . The maximum of metallic shell activity
shown in the ﬁgure can be estimated to be about T = 7300
(2015 – 2016). The maximum of equivalent width of the Hβline also has the same epoch as the metal lines. This is similar to the estimate from the I c variation of the Hα emission
4

5

Fujii Bisei Observatory⟨http://otobs.sakura.ne.jp/FBO/index.html⟩.
Since 1996, the observatory has published a huge amount of low dispersion spectra (R ∼ 500) focused on variable stars. The accumulation ofobservational data by ﬁxed spectroscopic instruments is attractive.
58th ﬁgure from the top in⟨http://otobs.sakura.ne.jp/FBO/fko/gcas/bu tau/bu tau.htm⟩

Vol. 6,

line, but it conﬁrms that the timing of the maximum shell activity in the group of metallic absorption lines formed at the
inner region of the disk is delayed by about one year from theHα line. Compared to Hirata’s (1995) estimate of the maximum shell activity period 1981 based on the Fe II absorption
equivalent width, the maximum period occurred about 34 years
later. Again, the periodicity of the disk formation activity issuggested.
3.2.2. I c variation after T = 9000
Figure 1 (bottom) is an enlarged view from T = 8300. After
T = 9000, the state of the I c short-term variations changes isvery different from before. A different I c variation appears
than before. Furthermore, the stable period after short-term
variations is not clear, and the I c value always increases during its period. The thin black line superimposed on the ﬁgure
is a curve tracing the I c ﬂuctuations, and the purple dashedline is a straight line connecting the observation data during
the period when the short-term variations were thought to ”disappear”. In the ﬁgure, we have tried to highlight the featurethat the pattern of the ”long-term background I c variation”
changes after T = 9000. Around T = 10000, I c shows always
I c > 1. In early March 2023, corresponding to the right end ofﬁgure 1, according to the Hβ observation (echelle, R = 9000)
in the BeSS database, the central absorption intensity of the Hβline increased and approached considerably the absorption intensity of the photospheric component. This is the time when
the shell absorption component has became very weak. Suchan I c variation is an interesting observation, ﬁrst observed at
the turning point of the repeated phase transition of Pleione, theprocess of changing from the Be-shell phase to the Be phase.
The Fujii Bisei Observatory has also published the long-term
spectral variation of Pleione for the region λλ 4200 – 5600 Å,
normalised to the continuous light6 . The two observations(January 2020 and September 2020) corresponding to the period around 9000 (just before the pattern of the long-term I cvariation changes) show that the metallic absorption lines are
already weakening.
These changes in the Hα, Hβ, and metallic lines are very
close to the observational guidelines for the onset of the Be
phase of Pleione (see ﬁgure 1 of Rivinius et al. 2013). The
March 2023 phase of Pleione is judged to be just before thephase transition to the Be phase. This is 35 years after thephase transition in the previous disk formation activity cycle(1988).
3.3. Radial velocity parameters (V r B, V r R, V r c)
Figure 2 (top) shows the evolution of the changes in the radial velocity parameters, the blue and red side emission peaks(V r B, V r R), and the central depth minimum (V r c) for allobservations. The values are converted to heliocentric radial
velocities, not corrected for the orbital motion. Figure 2 (bottom) is an enlarged view after T = 8300. As in ﬁgure 1, thepattern of variation can be seen as the superposition of theshort-term variations corresponding to periastron on the longterm background variation. We estimated the long-term background variation of V r c, using the same method as for the
6

57th ﬁgure from the top in⟨http://otobs.sakura.ne.jp/FBO/fko/gcas/bu tau/bu tau.htm⟩

Long-term evolution of the Hα line of Pleione

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Fig. 2. Long-term variation of the radial velocities of the Hα emission line, the blue emission peak V r B, the red emission peak V r R, and the central
minimum V r c. The notation is the same as in ﬁgure 1, except that the vertical axis is the heliocentic radial velocity for the three velocities. In the
V r c sequence the superimposed blue dashed line shows the long-term background variation, which is derived from the data outside the periastronphenomenon. (Top) All data from January 2009 to March 2023. (Bottom) Enlargement after July 2018. After T = 9000 the pattern of the long-term
background variation is tentatively highlighted by the dashed purple line.

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J. Katahira

I c variation. The result of the quadratic function ﬁt (correlation coefﬁcient is 0.53) is superimposed on the ﬁgure with ablue dashed line. According to Nemravavo et al. (2010) ﬁgure 10, the radial velocity of the primary star’s orbital motion
can be regarded as almost constant during the period outsidethese short-term variations. The blue dashed curve reaches amaximum value of 15.8 km s−1 at about T = 7800, which isan increase of about 7 km s−1 compared to T = 4800. The
long-term background V r c variation, which had a decreasing
trend after the maximum, begins to show a new trend of increasing V r c values after T = 9000 towards T = 10000, when
the spectral phase begins to shift to the Be phase (ﬁgure 2(bottom), ﬁgure 3 (bottom)). The mean values of V r c (kms−1 ) in the long-term background variation at some epochs
around T = 8500, 9180, 9840, and 10000 are 15.6 ± 1.6(37),16.5 ± 4.6(18), 16.8 ± 3.5(10), and 18.5 ± 4.1(4). The error
is the standard deviation and the number in parentheses is thenumber of observations. Although this change is within therange of observational error, it is likely to be related to the trend
in I c background variation in ﬁgure 1(bottom). Nemravova etal. (2010) ﬁgure 5 implies an increasing trend of V r c in the
spectral phase transition from Be-shell to Be of theprevious disk formation activity.
The emission peak separation PS (= V r R - V r B), readfrom the long-term background variation of V r B and V r R,
gradually narrows even during the maximum period of shell
activity. This change in PS is used to interpret the change indisk radius in ”4. Discussion”. However, the interpretation isbased on a rough approximation that considers Pleione to be asingle Be star once the periastron phenomenon of the companion subsides. During the periastron phenomenon the PS value
shows a tendency to decrease rapidly, but the physical picture
is not understood.
Short-term variations in radial velocity during periastronphenomena tend to be easy to read. First, V r B varies
in an increasing direction, and V r R in a decreasing direction, and then returns to its original state. The variation ofV r c seems to be related to the mean emission peak velocity
V r EP (= 0.5 × (V r B + V r R)), i.e. it seems to be affectedby the shift of the two side emission peaks (see ﬁgure 3). The
binary orbital motion causes the radial velocity of the primarystar to decrease by about 10 km s−1 before periastron and thento increase again (Nemravova et al. 2010). A plot corrected
for orbital motion would also be necessary, but the trend of theﬁgure does not change signiﬁcantly.
3.4. V/R and radial velocity variation
Figure 3 (top) shows variations for V/R (= I B/I R),
V r EP (= 0.5 × (V r B + V r R)) and V r c. Figure 3 (bottom) shows an enlarged view after T = 8300. To make it easierto read the long-term background variation of V/R, we haveﬁtted the data from the period when the periastron phenomena did not occur using the least squares method. The resultis shown as dashed lines (correlation coefﬁcient 0.86). The
value of the V/R in the background variation changes slowly inthe order of ”∼ 1”, ”> 1”, ”∼ 1”, and ”< 1”. The maximum
period of the V/R long-term background variation is about
T = 7800, about two years after the maximum period of Hαshell activity (about T = 7000) shown in ﬁgure 1. Note that asVol. 6,the Hα V/R reaches its maximum value, the Hβ line begins toshow V/R < 1 (Katahira et al. 2020), and as of March 2023,
V/R ≪ 1 (BeSS echelle spectrum).
To show the long-term background variation of the radial
velocity, we use the long-term background variation of V r c,which is indicated by the dashed line and is already plotted
in ﬁgure 2. Since the long-term background variations for the
V/R and the radial velocity follow a similar pattern between
T = 4800 and T = 9000, the correlation between them will be
discussed later (see ﬁgure 5). Note that in the backgroundvariation after T = 9000 in ﬁgure 3 (bottom), V r EP showsa similar variation to the V/R (purple dashed line), but V r c isclearly different. There is a difference in the pattern of variation between V r EP and V r c in the background variation.
The V/R short-term variations that occur during the periastron phenomena show changes that are roughly similar to theradial velocity variations. However, the phase of the radial velocity change is slightly behind the V/R short-term variation,
so the correlation does not appear. Although the Hα proﬁlevariation diagram is not shown in this paper, the time variationof the proﬁle shows a strangely ”regular” sequential intensitychange in the line during each periastron phenomenon . A detailed study of the shape change of the time series proﬁle isnecessary to understand the phenomenon.
4. Discussion
4.1. Interpretation of the long-term background variation ofIc
We wish to consider qualitatively the signiﬁcance of thelong-term background variation of the I c value from the image of the Be disk model. Rivinius, Carcioﬁ, and Martayan(2013) ﬁgure 1 shows the currently accepted disk model7 .
The change in proﬁle shape due to the direction of the observer’s line of sight (disk tilt) is shown with the proﬁles ofa typical Be star. From such a disk model, in the Be-shell
phase, Pleione is seen to be in the edge-on view. The particular state of Hα central minimum intensity I c < 1 in the Beshell phase of Pleione is considered to be due to the high density of the disk (e.g. Hummel 1994; Iwamatsu & Hirata 2008;
Silaj et al. 2014). Thus, the long-term background variation of
I c in ﬁgure 1 is thought to be indicative of the long-term variation that the disk exhibits around the edge-on state. We canconsider two factors that determine the long-term backgroundvariation of I c: disk density and inclination. It is a simpliﬁcation to say that what causes the almost symmetrical variation inﬁgure 1 is an increase or decrease in disk density, or a change
in disk inclination.
Let’s ﬁrst take a look at the plots of the measurements relating to the disk density parameter. The long term backgroundI c variation shown in ﬁgure 1 (top) at T = 4800 – 9000 showsa symmetric trend with a minimum I c = 0.22. Meanwhile,the emission line intensity is increasing. In ﬁgure 2 (top),
the long-term variation of the emission peak separation
7

The Be disk is currently considered to be, (i) geometrically, the disk has a
ﬂat shape that spreads out on either side of the equatorial plane of the disk
as one moves away from the central star, and (ii) physically, the disk is
a viscous decretion disk formed by the supply of the angular momentum
from the central star (Rivinius et al. 2013).

id. 1Long-term evolution of the Hα line of Pleione

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Fig. 3. Long-term variations of Hα emission peak intensity ratio V/R, central minimum velocity V r c and mean emission peak velocity
V r EP (= (V r B + V r R)/2). The notation is the same as in ﬁgure 1, except that the left vertical axis is V/R, and the right vertical axis is the heliocentric radial velocity for V r c and V r EP . In the V/R sequence the superimposed dashed line shows the long-term background variation, which isderived from the data outside the periastron phenomenon. The same is done in the V r c sequence. In the V r c sequence, the data outside the periastronphenomenon are represented by blue squares. (Top) All data from January 2009 to March 2023. (Bottom) Enlargement after July 2018. After T = 9000,
the patterns of the long-term background variations for the V/R and V r c sequences are tentatively shown by the purple dashed lines.

J. Katahira

4.2. Intensiﬁcation of the periastron phenomenon
The ﬁrst conﬁrmation of the periastron phenomenon seen
in ﬁgure 1 is at the periastron time of T = 4874, according tothe overlap of the periastron time and the ﬂuctuation in which
the I c value decreases in the short-term variation. The next
one is observed at T = 5092. After that, the short-term jumpof the I c value (correlated with the change of the emission
peak intensity) can be used as a guide, making it easy to judge
the periastron phenomenon. Looking at the ”magnitude” ofthe jump, based on the curve of the long-term I c backgroundvariation, it increases gradually towards the maximum periodof shell activity and increases exponentially after the maximum
period. This change clearly reﬂects the strengthening of the
disk-companion gravitational interaction at periastron, as thedisk radius increases and disk precession progresses.
Let’s imagine a process where the interaction strengthensat an accelerating rate towards the transition to the Be phase.
After T = 8500 in the middle, the narrowing of the emission
peak separation PS in the long-term background variation seen
in ﬁgure 2 is accelerated, and the disk radius is expected toincrease rapidly. The increase in the disk radius seems to be
the main reason for the increased interaction with the compan-

Ob

Hirata(2007)
Fig.3b

se
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62 deg.
T =10000

PS (= V r R - V r B) shows a decreasing trend in the corresponding period. For a single Be star, a simple relationshipbetween PS, disk inclination, and disk radius is considered (e.g.
Hirata & Kogure 1977). It is an approximation that, for thesame inclination, the disk radius increases as the PS decreases.
Even taking into account the change in the inclination due tothe precession of the disk, which will be discussed in the next
section, it can be seen that the size of Pleione’s disk continues
to increase. Since the increase in the size of the disk can beinterpreted as a result of the supply of gas from the central starto the disk, it is assumed that the density of the disk increases
during this time. It is then difﬁcult to generate a long term
background I c variation from up and down variations in thedisk density.
The interpretation of the long-term background variation of
I c will be explained in the remaining ”disk tilt”. Here weuse the image of the disk precession obtained by Hirata (2007)
from the long-term polarization analysis of Pleione. Since theobservations in the current disk formation period show similar changes to the previous disk formation period, so we assume that the same disk precession as in the previous period
is occurring. Focusing the disk arrangement near edge-on,
Hirata (2007) ﬁgure 3b, means that the disk gradually reaches
the state of inclination i = 90◦ (at the edge-on) and then separates again to reach the state of i = 60◦ (at the early stage of theBe phase). This picture of the precession can roughly explain
the present observation that the long-term background variation of I c in ﬁgure 1 gradually reaches a minimum and thenslowly increases again. Conversely, the long-term backgroundvariation of I c in ﬁgure 1 can be used to claim that the precession of the disk has ”reappeared”.
As an extension of the above reasoning, the observation thatthe value of I c is always I c > 1 at about T = 10000 may indicate that the disk inclination is approaching i = 60◦ . It isthought that the Hα emission line proﬁle for the Be phase begins to appear.

Vol. 6,
Inclination

8

1981.0 yr

2015.8 yr(FeII shell Maximum)(Metallic shell Maximum)

Time

Fig. 4. Approximation for calculating the current disk inclination

ion star. When the disk radius is calculated using the variation pattern of the disk inclination from Hirata (2007), theincrease in the radius at T = 10000 is about 1.7 times larger
than that at T = 8500. It is quite possible that the increasein the disk radius contributes to the intensiﬁcation of the periastron phenomenon. The present use of the inclination from
Hirata (2007) is based on the time of the previous and current
metallic shell activity maximum (1981.0, 2015.8 respectively).
The converted inclination values are i = 74◦ for T = 8500 and
i = 62◦ for T = 10000. This approximation process is illustrated by the hand-drawn ﬁgure (ﬁgure 4). This calculation
also suggests that at the beginning of the Be phase, the current
size of the disk may be larger than that of the previous disk
formation activity (see Hirata 2007 ﬁgure 3c).
Since the periastron phenomenon begins to intensify after
T = 9000 in spite of the gradual progress of the inclination
change, the factor caused by the change in the 3-dimensional
space arrangement may also be important. So let’s consider thefactors in the intensiﬁcation due to the arrangement of the disk,
whose inclination is approaching 60◦ due to the precession.
In the Hirata (2007) disk precession, the position of the diskchanges as if it were perpendicular to the orbital plane of thecompanion star. And in the initial state of the Be phase, thetilt of the disk reaches i = 60◦ . One of the factors that can be
deduced from such a spatial conﬁguration of the disk is thatthe light from the entire equatorial plane of the disk, where
the periastron phenomenon can occur, reaches the observer
more easily. In the period corresponding to the initial stateof Be phase(January – February 1988), Doazan, Thomas, andBourdonneau (1988) detected the appearance of superionizedlines of C IV and Si IV, while the Fe II shell absorption components weakened in the UV. Hirata (2007) interpreted this detection of UV superionized lines as the appearance of the ”stellar
wind bellow and above the cool disk” due to the disk precession. Because of the change in inclination, we naturally assume that the disk has begun to show new, previously unseen
regions, not only above the front face of the photosphere, butalso around the photosphere. It is very possible that such asituation has already appeared by March 2023 during the current disk formation period. The intensiﬁcation of the periastronphenomenon in question at around T = 10000 is thought to be
an integral effect of ”new areas” of the disk.
Another factor can be imagined as follows. Due to the precession, a change in the spatial conﬁguration of the disk in
the orbital plane of the companion star can create a situation

Long-term evolution of the Hα line of Pleionethat intensiﬁes the periastron phenomenon itself. This is because the gravitational interaction has an aspect that becomesstronger as the distance between the companion star and thedisk decreases near the periastron. The sharp change in the I c
ﬂuctuation shown on the right hand side of ﬁgure 1 could be
the result of the change in the spatial conﬁguration of the disk.
Elucidating the factors involved in the intensiﬁcation will clarify the interaction between the companion star and the disk.
During the intensiﬁcation of the periastron phenomenon, we
assume that the gas supply from the photosphere to the disk
continues. The observation of 13 February 2020 in the BeSS
database (T = 8892, R = 5700) shows that the broad and shell
components of the Ca II K line weaken considerably. The
intensity of this line is usually taken as a parameter indicating the degree of gas supply from the photosphere to the disk
(Gulliver 1977). However, we believe that the weakening tendency of the Ca II K line corresponds to the weakening of themetallic shell absorption lines introduced in section 3.2.
4.3. Traces of the V/R phenomenon in the long-term background variation
The V/R phenomenon is the periodic V/R variation of theemission line proﬁle that occurs in about 70% of Be-typestars (Mennickent & Vogt 1991). In the spectral type distribution, the number of stars decreases in the late type ofB8 type, such as Pleione (Okazaki 1997). Pleione has also
shown V/R phenomena during earlier periods of disk activity (Okazaki 1997; Hirata 2007). It has been shown theoretically that when the V/R phenomenon is observed, a stable onearmed density wave is activated in the Be disk and rotates inthe disk (Okazaki 1996). If the intensity of the emission peakon the blue side of the proﬁle is stronger than the intensity ofthe emission peak on the red side (V/R > 1), the whole proﬁleshifts to the red side, and vice versa. Let’s consider whether the
V/R phenomenon, which shows such a correlation, also occurs
during the current period of disk formation activity.
From ﬁgure 3 it can be seen that the long-term backgroundvariations of V/R and V r EP follow a similar pattern. Since
the value of V r EP can be considered as roughly corresponding to the representative wavelength position of the proﬁle, thecorrelation between the two parameters is plotted in ﬁgure 5 tocheck the existence of the V/R phenomenon. The marks plotted are data up to T = 9000 for the black marks, and data above
than T = 9000 for the red marks (the observation range for the
background variation is not clearly found). We have the impression that the two sets are related over the whole V/R range.
The ﬁgure means that the whole emission line shifts to the redside when V/R > 1, but starts to shift to the blue side when
V/R = 1. Therefore, when all the data are ﬁtted with a straightline using the method of least squares, a blue dashed line is obtained (correlation coefﬁcient 0.53). A positive correlation is
tentatively observed between V/R and V r EP .
For the ﬁrst time since Pollmann and Vollmann (2014)
pointed to the existence of the periastron phenomenon, we
have identiﬁed traces of the so-called V/R phenomenon in thelong-term background variation, a series of discrete data forthe period outside the periastron phenomenon. However, after
T = 9000 the selected data include data that cannot be visuallyjudged to be in a state where the periastron phenomenon has

9Velocity(km/s)

id. 1

Vr_EP
4852<T<9000
9000<T<10030

20

Corr. Coeff. = 0.53

0

0.9

1

V/R

1.1

Fig. 5. Correlation diagram of V/R and V r EP . The data are for observation days considered to be out of the periasron phenomenon, andare color-coded by T = 9000.

subsided. Data with V/R < 1 are expected to be observed in
the future, but it may be necessary to change the method of
selecting data periods representing the long-term backgroundvariation from using I c intensity to using emission line peakintensity.
Note that the correlation between V/R and V r(I c) shows ascatter similar to the black squares in ﬁgure 5 up to the 9000
timescale, but the correlation disappears after T = 9000. The
reason for this is that the variation pattern of the V r(I c) isdifferent from that of V/R after T = 9000, as shown in ﬁgure 3(bottom).
4.4. Can the V/R phenomenon in the long-term backgroundvariation and the short-term V/R variation in the periastron coexist?
As can be seen in ﬁgure 1, the disk is gravitationally shaken
with each periastron passage of the companion star. The V/Rphenomenon in the long-term background variation shown insection 4.3 is only conﬁrmed by connecting the data observed
during the period when the periastron phenomenon seems tohave subsided, so there is some uncertainty. However, let
us dare to assume that the correlation diagram in ﬁgure 5
shows the appearance of the V/R phenomenon. Then thelong-term background V/R phenomenon in the disk coexists
with the short-term V/R variation during the periastron phenomenon. The variation of the V/R value in the periastronphenomenon shown in ﬁgure 3 is about twice as large as
the variation considered in the long-term background variation in ﬁgure 5. Furthermore, in the Hβ line formed in the
inner region of the disk, the V/R variation synchronised with
the Hα line is observed during the periastron phenomenon(Katahira et al. 2020). This indicates that the gravitational inﬂuence of the companion extends into the inner region of the
disk. The existence of the V/R phenomenon in the Hα line
means that the density wave is maintained throughout the disk,even though the periastron phenomenon occurs repeatedly andaffects deep within the disk. It will be important to clarify
whether the periastron phenomenon in the short-term variationand the V/R phenomenon in the long-term background varia-

10

J. Katahiration really coexist.
Panoglou et al. (2018) simulated the Hα line proﬁle variations (V/R variations) produced by a double spiral density wave
in the Be disk of a close binary system. Interestingly, they show
something similar to the V/R variations that appear during theperiastron phenomenon. This could be a clue to thinking about
the periastron phenomenon.
5. Summary
The Hα data of Pleione in the BeSS database are collective observations by amateur spectroscopic observers, providing the valuable high temporal resolution record. This paper
uses 520 observations (wavelength resolution 10000 – 20000)
in the database from January 2009 to March 2023. From theHα emission proﬁle normalised to the continuous light level,
we measured the intensities of the blue and red side emission
peaks, the radial velocities at the peak intensities, the intensityof the central minimum between two emission peaks, and theradial velocity at the minimum. We have plotted the 14-year
variation of the measured data for the above three characteristic points of the proﬁle, and discussed the following:• The 14-year variation is tentatively divided into the repeated short-term variations and the long-term background variation.
• The locus of the long-term background variation of thecentral minimum intensity may be a reﬂection of the disk
precession.
• The V/R phenomenon may be occurring in the long-term
background variation.

Ackowledgment
This work has made use of the BeSS database, operated at
LESIA, Observatoire de Meudon, France, and the results of
the long-term analysis of Pleione published by the Fujii Bisei
Observatory. The author wishes to express his hearty thanks
to the BeSS database and the Fujii Bisei Observatory. The author is grateful to the referee for the useful comments, which
helped to improve the manuscript. Stimulated by the proofof the existence of the periastron phenomenon by Pollmann
and Vollmann (2014), I started to study the data of Pleione in
the BeSS database with the help of Drs. R. Hirata (ex-Kyoto
University) and S. Honda (University of Hyogo). I learned agreat deal from the discussions with both of them at that time.
I am deeply grateful to both of them.
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