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Up: Long-term visual spectrophotometric

3. Results

The spectrophotometric data gathered for all program stars are presented in Fig. 6 (available in electronic form). For each star there are three panels. The first shows the observed time variability of the V, tex2html_wrap_inline4779 and D. The second panel shows the (U,D), (V,D) and (tex2html_wrap_inline4789) correlations and the third the (tex2html_wrap_inline4791) correlation. The symbols used are: black dots for the BCD data, crosses for UBV data, stars for UBVRI, triangles for Geneva photometry, diamonds for the uvby system and squares for 13-colour data. The following comments for individual stars are divided into two subsections according to the degree of known spectrophotometric and spectroscopic variations.

3.1. Spectrophotometric and spectroscopic behaviour of some well-studied Be stars

This subsection deals with program Be stars, for which we have enough BCD spectrophotometric observations to allow us: (a) to estimate the uncertainties discussed in Sect. 2.6; (b) to assign not only a degree of reliability to relations derived using only indirect (tex2html_wrap_inline5887) parameters, but also to judge the likelihood of scatters and more or less systematic deviations from the sketched mean relations. As for all these objects rather copious spectroscopic observations exist, in this section we also summarize their main spectroscopic variations over the last fifty years, more or less simultaneous to photometric data which are analyzed in this paper. The qualitative description of the correlated spectroscopic and spectrophotometric behaviours of these objects can be used as a reference for the remaining program stars, for which only spectrophotometric variations are presented. These stars are divided into two groups: (1) Be stars most commonly seen in Be phases: tex2html_wrap_inline5089 Cas, X Per, 11 Cam, tex2html_wrap_inline5293 Oph and 59 Cyg; (2) Be stars most frequently observed in Be-shell phases: Pleione, 48 Lib and 88 Her.

3.1.1. Stars most frequently seen in Be phases

tex2html_wrap_inline5089 Cas (HD 5394)

Observed since 1866, tex2html_wrap_inline5089 Cas firstly showed a long Be phase of strong emission. Then, between 1932 and 1942, dramatic changes occurred in its spectrum; two strong Be phases, each followed by a strong shell phase, have been observed (Doazan et al. 1983 and references therein). In 1942, tex2html_wrap_inline5089 Cas entered a quasi-normal B phase. Four years later, a subsequent Be phase increased slowly and irregularly. Weak secondary minima in emission were reported in 1956-1961 (Hubert-Delplace & Hubert 1979 hereafter referred to as "Atlas''), end 1975, 1983-1984 (Horaguchi et al. 1994 and references therein) and in 1990-1991 (Peters 1990, 1991); a "veiling'' effect of lines was seen from the end of 1963 to 1972 (Atlas). According to Horaguchi et al. (1994), Htex2html_wrap_inline5899 intensity has oscillated with a time scale of several years. Long-term far UV and visual variability seems to be associated (Doazan et al. 1987).

It is interesting to note that the minima in the tex2html_wrap_inline4779 gradient correspond to minima in emission strength and the maximum of tex2html_wrap_inline4779 in 1966 to a veiling in line spectra. Furthermore the total BD is strongly negative at the epochs of strong outbursts (1932-1942). Somewhat scattered but single (tex2html_wrap_inline4779, D) and (tex2html_wrap_inline4779, V) correlations with tex2html_wrap_inline5915 slopes are present, although photometric and spectrophotometric data correspond to epochs before and after a number of Be tex2html_wrap_inline4885 B-normal tex2html_wrap_inline4885 Be-shell phase transitions. We note, however, that the (V, D) relation has two branches in tex2html_wrap_inline5925 dex. Let us finally note that the rapid photometric variability which is commonly associated with surface stellar activity is still debated for this object.

X Per (HD 24534)

Variable but generally weak emission was noted in the spectrum of X Per before the end of the forties (Cowley et al. 1972). In 1951-1953, a strong emission phase phase with a "veiling'' effect began; afterwards marked strengthening (outburst) of emission lines (HI, HeI, FeII, SiII) and of the "veiling'' effect from 1957 to 1961 has been reported, with a maximum in 1961 (Cowley et al. 1972; Wackerling 1972, Atlas). In 1974 the emission decreased for 2 years, becoming bright again by the end of 1976. In 1978, only very weak emission was seen in Htex2html_wrap_inline5899, the spectrum was consistent with a quasi-normal O9-B0 star (de Loore et al. 1979). By the end of 1978, a new emission phase reappeared (Roche et al. 1993 and references therein). Still fairly bright in February 1990, the Htex2html_wrap_inline5899 line was seen converted to absorption in September 1990. A strong shell feature superposed on the photospheric component of Htex2html_wrap_inline5899 and HeI6678 in Nov., Dec. 1990 and Jan. 1991 was reported by Norton et al. 1991. Still in absorption in Sept. 1991 (Reynolds et al. 1992), weak emission was again present on Htex2html_wrap_inline5899 in Oct. 1991 (Kaper & van Kerkwijk 1992).

Decrease and minima of line emission generally correspond to lower brightness states. Single (tex2html_wrap_inline4779, D) and (tex2html_wrap_inline4779, V) correlations are accompanied by bifurcated or double-fold (U, D), (V, D) relations with tex2html_wrap_inline5951 slopes in D < D* corresponding to definite Be phases. The stellar BD is D* = 0.050 dex (de Loore et al. 1979). For tex2html_wrap_inline5959 incipient horizontal relations: tex2html_wrap_inline5961 are apparent which correspond to the shell phase.

11 Cam (HD 32343)

11 Cam has always had strong emission observed up to high Balmer lines. Often observed before 1975, this star has presented recurrent outbursts. Schild (1973) reported outbursts in 1916-1919, 1929-1931 and 1943-1949. From 1953 to 1974, emission was always strong in Balmer lines with a weak minimum in 1961. Bright and particularly clear FeII lines were reported from 1963 to 1971 during maximum emission. Afterwards they seem to have disappeared (Atlas). Isolated observations between 1980 and 1989 of the Htex2html_wrap_inline5899 emission line have revealed a rather constant strength of this line except in October 1983 (Ballereau & Chauville 1987) for which the line is very strong (no photometry available at that time). Conversely no spectra are available at maximum brightness in November 1977. The rapid photometric variability of this star is suspected although not yet confirmed.

Single (U, D), (V, D), and (tex2html_wrap_inline4779, V) relations are seen, although they all have tex2html_wrap_inline5977 and tex2html_wrap_inline5979 slopes. The (tex2html_wrap_inline4779, D) relation is scattered and single, however it has a tex2html_wrap_inline5985 slope.

tex2html_wrap_inline5293 Oph (HD 148184)

Since 1919 tex2html_wrap_inline5293 Oph has been characterized by strong emission in Balmer lines (Schild 1973; Slettebak 1982). A very strong emission in Balmer and FeII lines between 1953 and 1973 with a veiling effect seen up to 1969 (Atlas). From 1979, FeII emission lines have decreased in intensity; they were barely visible in 1975. Andrillat & Fehrenbach (1982) and Dachs et al. (1981) noted a steady decrease in the strength of Htex2html_wrap_inline5899 emission from 1972 to 1980. Then it seems to have slightly increased in 1982-1983, with a maximum through March 1983 (Dachs et al. 1986), and afterwards fluctuated with maxima tex2html_wrap_inline5993 in 1985, 1989 and 1993 (Hanuschik et al. 1996). The decrease of the Htex2html_wrap_inline5899 emission line from 1972 to 1980 corresponds to an increase of brightness and to reddening of the gradient tex2html_wrap_inline4779. Total Balmer discontinuity D varied strongly during this period with a strong minimum near 1983. D is at minimum when tex2html_wrap_inline6003 of Htex2html_wrap_inline5899 emission is at maximum. tex2html_wrap_inline5293 Oph is a rapid photometric variable star (Balona 1995) with a moderate tex2html_wrap_inline4793 (< 130 km s-1).

The (U, D) and (V, D) relations are single for D <0 with tex2html_wrap_inline5977 slopes, but two branches with tex2html_wrap_inline5951 slopes seem to appear at tex2html_wrap_inline6029. The (tex2html_wrap_inline4779, D) and (tex2html_wrap_inline4779, V) relations are scattered and whilst the first is likely to be single with a tex2html_wrap_inline5985 slope, the second seems to be double valued with one of the branches with a tex2html_wrap_inline5979 slope.

59 Cyg (HD 200120)

Since 1904, 59 Cyg has presented active and quiescent phases in the behaviour of its CE, successively producing spectra with the appearance of a quasi-normal B, Be and Be-shell star. Long-term, rather stable, strong emission-line phases alternate with phases of marked changes, generally characterized by slow increase and shorter decreases of the strength of emission lines. A slow increase of emission was observed from 1945 to 1950 (Merrill & Burwell 1949). Afterwards, a quiescent phase set in from 1953 to 1970, the emission features being relatively stable; emission was at maximum in 1956 and in 1961, and at minimum in 1967 (Atlas). Short-lived outbursts were then observed from July to the end of 1972 and from December 1973 to June 1974, each followed by a strong shell phase of several months duration. After that, emission vanished and almost completely disappeared in 1976-1977 (Hubert-Delplace & Hubert 1981). An increase of emission started in the end of 1979 and ended near mid-1980. Emission decreased again to a minimum in 1982, and again increased from 1983 to 1986. From 1979 to 1985 strong Htex2html_wrap_inline5899 emission strength changes occurred, being correlated with the wind observed in the far UV (Doazan et al. 1989). From 1986, a new, rather stable quiescent emission phase has been established (Peters 1989-1992, 1994).

Strong emission lines are present when D < D* and the line emission decrease corresponds to a progressive increase of the BD, which ended at D = D* in 1977 when the line emission almost disappeared. From the existing photometric data, it is seen that after 1961 the mean V magnitude suddenly began decreasing until 1977, fluctuating at intervals of tex2html_wrap_inline6053 yrs. with a variable, damped-like amplitude whose maximum was nearly 0.4 mag. The intrinsic continuum reddening seems to be strongest when the star is brightest but not necessarily well correlated with line emission intensity as tex2html_wrap_inline4779 remains high even in 1967 at a relative minimum of line emission. This reddening cleared up in 1977, so that tex2html_wrap_inline4779 became that of a B1V normal star for D* = 0.116 dex shown by a dashed vertical line in Fig. 6. Unfortunately, there are no photometric data corresponding to epochs of short-lived shell phases in 1973 and 1974. It is seen, however, that the (U,D), (V,D), (tex2html_wrap_inline4789) and (tex2html_wrap_inline4791) relations are single with tex2html_wrap_inline5951 and tex2html_wrap_inline5915 slopes and that they are nicely conserved, even after the shell phases and the periods of loss of emission characteristics. A photometric microvariability of this star was observed, but there is no period determination at the present time.

3.1.2. Stars most frequently seen in Be-shell phases

Pleione (HD 23862)

Notable long-term changes are reported in the literature, for this star, see Hirata (1995) and references therein. Recently this star was discovered by Katahira et al. (1996) to be a binary with an orbital period of 218.0 days. We give here only a summary of spectral variations. Observed in a Be phase until 1904, Pleione entered a phase showing a rapidly rotating B normal-like star from 1905 to 1936. The first Be-metallic shell phase occurred in 1938 and ended in 1954. From 1954 to 1972, this star again entered a Be phase with strong Balmer and FeII emission lines at maximum around 1960. Then gradual weakening began; in 1972, the Htex2html_wrap_inline5899 emission line was very weak. In late 1972 another metallic shell phase started and developed with a maximum near mid-1982. The metallic lines faded and disappeared in 1988 (Hirata 1995). The star entered a new Be phase; the Htex2html_wrap_inline5899 emission line gradually increased from 1982 (epoch of maximum strength in shell lines) to 1993 and decreased after. The global behaviour of spectral and light changes has been studied by many authors (Hirata & Kogure 1976; Kogure & Hirata 1982; Goraya & Tur 1988). A steep decline in brightness occurred in 1938 and 1972 as each respective shell phase started. Furthermore, a strong weakening in the U magnitude came with the maximum of metallic shell lines in 1982-1983.

From our study it is shown that in 1972, at the epoch of the appearance of a new Be-metallic shell phase associated with a rapid fainting in the UBV colours, the tex2html_wrap_inline4779 gradient rapidly increased and total BD decreased indicating the contribution of a secondary BD in emission. Between 1972 and 1988 the behaviour of the total BD was the same as the strength of the metallic shell lines. It is interesting to note that the maximum value of the total BD reported in 1982 corresponds to a maximum of the absorption strength of metallic lines. In the short interval of D < D*, this star shows two (U,D) and (V, D) relations which are continued in the region of D > D*, where normally in Be shell phases there is a single relation. The (tex2html_wrap_inline4779, D) and (tex2html_wrap_inline4779, V) seem to be single, unless for D < D* possible double relations are not well resolved. The (tex2html_wrap_inline4779, D) relation is slopeless, while (tex2html_wrap_inline4779, V) seems to have a global tex2html_wrap_inline5979 slope. Rapid photometric variability was found in this object (McNamara 1987).

48 Lib (HD 142983)

48 Lib is one of the best examples of shell stars. The variable shell episode of 48 Lib started in 1935: RV cycles (8-13 yr), marked variations of intensity and in line profiles. These phenomena become more striking with increase of the amplitude of radial velocity (RV) (Aydin & Faraggiana 1978 and references therein; Mon 1984; Hubert et al. 1987). In 48 Lib successive cycles exhibit negative RV phases longer than positive ones. The strength of shell lines has varied strongly since 1935. It was near or at maximum in 1943, 1965, 1970, 1974 and towards the end of 1993, according to recent observations published by Hanuschik et al. (1996). Maxima in the strength of emission in the first Balmer lines were generally seen at the same time or very close. Long-term photometric variations associated with spectroscopic cycles, and mid-term quasi-periodic oscillations (10-20 days) and sudden fadings (1-2 days) were reported by Mennickent et al. (1994) . These authors noted a brightness maximum in 1988 and possible minima in 1982 and 1991. In compiled data presented in this study, pronounced maxima of brightness were observed in 1976 and 1988. Balmer discontinuity is respectively lower and higher at epochs close to minima and maxima of the strength of shell metallic lines. Single relations are seen in (U, D), (V, D), (tex2html_wrap_inline4779, D) and (tex2html_wrap_inline4791). The (U, D) has a tex2html_wrap_inline6145 slope, while in all others there is no slope. This star is a rapid photometric variable, which may in part explain the width of spectrophotometric relations.

88 Her (HD 162732)

88 Her is well-known to have presented associated long-term changes in its spectrum and in its light; in addition it was discovered by Harmanec et al. (1972) to be a binary (P = 86.72 days). This star exhibited a Be phase with a hydrogen and metallic shell from 1955 to 1960. The Htex2html_wrap_inline5899 emission line gradually weakened from 1961 to 1971; the metallic shell lines, very weak from 1966, disappeared near 1970 (Atlas). In 1972-1976, the star was in a quasi-normal B phase, the Htex2html_wrap_inline5899 emission line was very weak in 1976-1977, and at the same epoch a minimum of the linear optical polarization was seen (Arsenijević et al. 1987). Then, the Htex2html_wrap_inline5899 emission line again gradually increased from 1977 to 1986. In addition, from 1978 to 1982, a new metallic shell phase was increasing, followed from 1983 to 1986 by a decreasing metallic shell and mild Be phase (Doazan et al. 1982a; Hirata 1995). It was reported that the gradual weakening of Balmer emission and metallic shell lines in 1966-1970 was associated with brightening of the star (Doazan et al. 1982b) and that a large and rapid drop of the luminosity occurred near 1978, just before the development (1978-1982) of a new Be-shell metallic phase (Harmanec et al. 1980; Doazan et al. 1986; Hirata 1995). The polarization percentage rapidly increased in 1978, followed by a slow decline from 1979 to 1985 (Arsenijević et al. 1987).

From this study, it is seen that the total BD was at minimum when the metallic shell lines disappeared towards 1970 and during the quasi-normal B phase (1972-1976). Then it gradually increased to mid-1982, in the same manner as the strength of shell lines. On the other hand, the star became brighter as the equivalent width of the Htex2html_wrap_inline5899 emission line increased (1978-1986), according to Hirata (1995). It is worth noting that BCD observations of the stellar BD from mid-1977 to the end of 1983 do not reveal any change in the photospheric tex2html_wrap_inline6163 and log g parameters of this star during its phase changes (Zorec et al. 1989). The difference of brightness in the U and V magnitudes between the regions of D < D* and D > D* are clearly seen in Fig. 6. Single and slopeless relations (U, D), (V, D), (tex2html_wrap_inline4779, D) and (tex2html_wrap_inline4779, V) are characteristic for D > D*. No rapid photometric variability was found in this object.

3.2. Brief remarks on the remaining Be stars

HD 28497. A rapid photometric variable star. The emission is strong when D is small. The emission was low from 1989 to 1992; since then, D and V are increasing.

HD 35439. After a B normal-like phase, Htex2html_wrap_inline5899 emission appeared when D was at its minimum value.

HD 37202. A shell and rapidly variable star. D is at minimum when the strength of the shell is at minimum, but in such a case the RV is not at minimum in this star.

HD 37967. Emission is strong and there is almost no variation.

HD 48917. A rapidly variable star. Only a few spectroscopic data exist for this star.

HD 50123. Interacting binary having a composite B spectrum with a giant K companion. "Ellipsoidal'' photometric variations with a period of P = 28.6 days were found by Sterken et al. (1994).

HD 56014. Star with a moderate rotation (tex2html_wrap_inline6207 km s-1) showing at epochs a shell behaviour D > D*.

HD 56139. A rapidly variable star with a small tex2html_wrap_inline4793 (60 km s-1) but sometimes showing the shell behaviour D > D*. The value of tex2html_wrap_inline6003 for Htex2html_wrap_inline5899 is rather high (tex2html_wrap_inline6223). When D is smallest, Htex2html_wrap_inline5899 emission is weak.

HD 58978. A helium-shell star. The strong spectroscopic variations in 1990 and in 1991 do not seem to be apparent for the spectrophotometric variations.

HD 60848. One of the hottest Be stars where the Balmer and Paschen line emission is well correlated with the Balmer continuum emission.

HD 60855. Star with moderate line emission showing B tex2html_wrap_inline6229 Be phase transitions. There are no strong spectrophotometric variations.

HD 63462. Star showing slight variations in the Htex2html_wrap_inline5899 emission intensity.

HD 65875. A rapidly variable star with moderate tex2html_wrap_inline4793 (150 km s-1) and showing however phases with the shell behaviour D > D*. The star has a high value of tex2html_wrap_inline6003 in Htex2html_wrap_inline5899 (tex2html_wrap_inline6243). Slopes tex2html_wrap_inline5977 and tex2html_wrap_inline5979 seem to exist for this star.

HD 68980. Star with a very moderate tex2html_wrap_inline4793 (95 km s-1) and with very strong Htex2html_wrap_inline5899 emission (tex2html_wrap_inline6255). D seems to be smallest when Htex2html_wrap_inline5899 emission is strongest. Two relations between U, V versus D and for tex2html_wrap_inline4779 against V are apparent.

HD 83953. A rapidly variable star where D is probably at maximum when the emission is low.

HD 89890. A star with a low rotation, tex2html_wrap_inline6273 km s-1, having D > D*.

HD 91120. Emission is moderate with few variations.

HD 109387. A rapidly variable star. Emission is strong when D is low and tex2html_wrap_inline4779 is high. There is a high dispersion in the (D,V) diagram for tex2html_wrap_inline6285 in 1986.

HD 120991. Star with small tex2html_wrap_inline4793 (70 km s-1). Sometimes tex2html_wrap_inline6003 is high for Htex2html_wrap_inline5899 (tex2html_wrap_inline6295 in 1993). D diminishes when the emission increases. There are probably two correlations for U and V as a function of D.

HD 131492. A temporary strong shell phase was observed by Slettebak in 1980 which corresponded to a minimum in V and to a maximum in D (Slettebak 1982). This is responsible for the horizontal part in the diagram of tex2html_wrap_inline4779 against D. The shell behaviour D > D* is clearly seen, although the star apparently has moderate rotation (tex2html_wrap_inline4793 = 100 km s-1).

HD 162428. The emission and the shell are variable.

HD 168797. This is a B tex2html_wrap_inline6229 Be phase variable star. Emission is moderate.

HD 171780. A B tex2html_wrap_inline6229 Be variable star with moderate emission.

HD 178175. Star observed in spectroscopy very irregularly. The shell behaviour D > D* is slightly present and the star has a moderate rotation (tex2html_wrap_inline4793 = 120 km s-1).

HD 183656. A rapidly variable shell star. D is correlated with the V/R variations (the V/R and radial velocity curves are given in Koubsky et al. 1989). D is at maximum when the RV is at maximum.

HD 184279. A shell star, where D is correlated with the magnitude V and with the RV curve: D is at maximum as RV is maximum (in 1979).

HD 187811. A B tex2html_wrap_inline6229 Be variable star with rather moderate emission.

HD 191610. A rapidly variable star. It has been seen in B phase in 1955 to 1958, and in a Be phase since then. Emission is moderate.

HD 195325. A late Be star with a hydrogen and metallic shell.

HD 205637. A rapidly variable star where the hydrogen and metallic-line shell are variable.

HD 217050. A rapidly variable shell star. D is at maximum when the strength of the shell is at maximum.

HD 217543. Star with moderate emission showing B tex2html_wrap_inline6229 Be phase transitions.

HD 218674. Star with rapid variability. It also shows strong emissions and a hydrogen-shell.

3.3. Comments on the correlations obtained

3.3.1. General spectrophotometric behaviour

We may conclude that patterns in Fig. 6 show that spectrophotometric changes in Be stars are characterized by tex2html_wrap_inline6349 and tex2html_wrap_inline6351 relations which differ on the emission/absorption phase and may differ from star to star. More or less well-defined relations involving D < D* correspond to definite Be-phases. They can be either single or bivalued, mostly with tex2html_wrap_inline5951 and tex2html_wrap_inline5985, but sometimes also with slopes tex2html_wrap_inline5951 and tex2html_wrap_inline6361. On the contrary, the most current shape of relation during SPh-shell phases where D > D* can be summarized by: tex2html_wrap_inline6365, and depending on the stars tex2html_wrap_inline6367 or tex2html_wrap_inline6145. The constancy of the magnitude V is within 0.15 mag and sometimes even more. However, for some Be-shell stars, variation in V can be as high as 0.25 mag.

In general, of those Be stars which have two correlations of U, V against D, and, tex2html_wrap_inline4779 against V in Be phases (D < D*), nearly all have strong Balmer emission lines. The fact that there are two relations or only one is not related to particular values of tex2html_wrap_inline4793.

3.3.2. Stars with low tex2html_wrap_inline4793 in SPh-shell phases

It is noteworthy the behaviour of some stars like HD 56014, HD 56139, HD 65875, HD 131492, where all have D* determined in the BCD system, which show both SPh-Be and SPh-shell phases. The mean deviation of these objects for the SPh-shell phase is tex2html_wrap_inline6395 dex which is more than 3 times the expected error tex2html_wrap_inline5739 given in Table 3. Most of these deviations are depicted by data from the Geneva photometric system, which is one of the most stable and uniform (Sterken & Manfroid 1991). These comments are also relevant to HD 178175, where D* is only for the mean MK spectral type and the tex2html_wrap_inline6401 dex deviations of the SPh-shell like phase are established from the Geneva photometric data. In Fig. 6 we see that for most of the above mentioned stars, the transition between a SPh-Be and a SPh-shell phase is characterized by a change of slopes in the spectrophotometric correlations. It would be difficult to understand such slope changes in terms of errors affecting the BDs, as they would likely conserve one of the observed slopes.

Among those stars where both spectrophotometric behaviours: Be (tex2html_wrap_inline4973; D < D*) and shell (tex2html_wrap_inline4879; D > D*) were seen, there are some with small or moderate tex2html_wrap_inline4793. We also note that two Be stars with low tex2html_wrap_inline4793 were seen only in SPh-shell phases where tex2html_wrap_inline6415 (HD 89890 (tex2html_wrap_inline6417 km s-1) and HD 178175 (tex2html_wrap_inline6421 km s-1). Low values of tex2html_wrap_inline4793 may correspond either to small tex2html_wrap_inline6427 or to low velocity V. In the first case, the SPh-shell phase: D > D*, cannot be explained by highly flattened CE seen pole-on. In the second case, we should admit the existence of an important fraction of slowly rotating Be stars (Mennickent et al. 1994). Nevertheless, for most Be-shell stars studied in this paper it is tex2html_wrap_inline6433 300 km s-1.

On the other hand, it is also worth noting that in Be-shell stars where enough RV data existed to be correlated with D (HD 142983, HD 183656, HD 184279 etc.), the highest RV > 0 appear when the line shell phenomenon and the absorption in the BD due to the CE are strongest. This phenomenon might favor the formation of compact CE layers near the star.

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