A&A 368, 912-931 (2001)
DOI: 10.1051/0004-6361:20000577

Statistical analysis of intrinsic polarization, IR excess and projected rotational velocity distributions of classical Be stars

R. V. Yudin^1,2

1 - Central Astronomical Observatory of the Russian Academy of Sciences at Pulkovo, 196140 Saint-Petersburg, Russia
2 - Isaac Newton Institute of Chile, St.-Petersburg Branch, Chile

Received 9 December 1999 / Accepted 22 December 2000

Abstract
We present the results of statistical analyses of a sample of 627 Be stars. The parameters of intrinsic polarization $(p_{\ast})$, projected rotational velocity $(v \sin{i})$, and near IR excesses have been investigated. The values of $p_{\ast }$ have been estimated for a much larger and more representative sample of Be stars ($\approx$490 objects) than previously. We have confirmed that most Be stars of early spectral type have statistically larger values of polarization and IR excesses in comparison with the late spectral type stars. It is found that the distributions of $\,p_{\ast}\,$ diverge considerably for the different spectral subgroups. In contrast to late spectral types (B5-B9.5), the distribution of $\,p_{\ast}\,$for B0-B2 stars does not peak at the value $p_{\ast}=0$%. Statistically significant differences in the mean projected rotational velocities ( $\overline{v\sin{i}}$) are found for different spectral subgroups of Be stars in the sense that late spectral type stars (V luminosity class) generally rotate faster than early types, in agreement with previously published results. This behaviour is, however, not obvious for the III-IV luminosity class stars. Nevertheless, the calculated values of the ratio $v_{\rm t}/v_{\rm c}$ of the true rotational velocity, $v_{\rm t}$, to the critical velocity for break-up, $v_{\rm c}$, is larger for late spectral type stars of all luminosity classes. Thus, late spectral type stars appear to rotate closer to their break-up rotational velocity. The distribution of near IR excesses for early spectral subgroups is bi-modal, the position of the second peak displaying a maximum value $E(V-L)\approx1\mbox{$\,.\!\!\!^{\rm m}$ }3$ for O-B1.5 stars, decreasing to $E(V-L)\approx0\mbox{$\,.\!\!\!^{\rm m}$ }8$ for intermediate spectral types (B3-B5). It is shown that bi-modality disappears for late spectral types (B6-B9.5). No correlations were found between $p_{\ast }$ and near IR excesses and between E(V-L) and $v\sin{i}$ for the different subgroups of Be stars. In contrast to near IR excesses, a relation between $p_{\ast }$and far IR excesses at 12 $\mu$m is clearly seen. A clear relation between $p_{\ast }$ and $v\sin{i}$(as well as between $p_{\ast }$ and $\overline{v\sin{i}}/v_{\rm c}$) is found by the fact that plots of these parameters are bounded by a "triangular" distribution of $p_{\ast }$:$v\sin{i}$, with a decrease of $p_{\ast }$towards very small and very large $v\sin{i}$ (and $\overline{v\sin{i}}/v_{\rm c}$) values. The latter behaviour can be understood in the context of a larger oblateness of circumstellar disks for the stars with a rapid rotation. From the analysis of correlations between different observational parameters we conclude that circumstellar envelopes for the majority of Be stars are optically thin disks with the range of the half-opening angle of $10\hbox{$^\circ$ }<\Theta<40\hbox{$^\circ$ }$.

Key words: classical Be stars: polarization - projected rotational velocities - near IR excesses - far IR excesses

1 Introduction

One of the challenging questions in Be stars investigations is the geometrical form of their circumstellar (CS) envelopes. There is still active debate as to the average value of the opening angle of CS disks around classical Be stars and on the application of the wind-compressed disk model (WCD). The ideas on this matter at present are rather controversial. Most authors have considered geometrically thin CS disks with half-opening angles $\Theta<3\hbox{$^\circ$ }$ as follows from the WCD theory of Bjorkman & Cassinelli (1993) and Owocki et al. (1994) (see for example Wood et al. 1997; Quirrenbach et al. 1997 etc.). However, acceptance of such a narrow disks faces several problems in the interpretation of spectroscopic characteristics (Moujtahid et al. 1999; Rivinius et al. 1999) and the observed IR excesses in Be stars (Porter 1997). Moreover, the observed phase transitions - Be $\Leftrightarrow$ Be-shell $\Leftrightarrow$ normal B - cannot be explained in the framework of a geometrically thin disk model (Moujtahid et al. 1999). Others used disk models with larger values of $\Theta\approx6\hbox{$^\circ$ }$ (Porter 1996) or $\Theta\approx13\hbox{$^\circ$ }\rightarrow15\hbox{$^\circ$ }$ (Hanuschik 1996; Waters et al. 1987) removing some of the difficulties. Finally, some authors recently returned to the model of flattened ellipsoids or spheroidal envelopes (Moujtahid et al. 1999) similar to those proposed earlier by Doazan & Thomas (1982). Although a low level of intrinsic polarization of Be stars is considered as an indication that the envelope must be strongly flattened, other geometries can produce the same level of polarization (see Wood et al. 1996a, 1996b). The IR excesses and spectroscopic characteristics, considered separately, also do not provide unique interpretation on CS geometries. Nevertheless, it is well established that the observed IR excess and polarization of radiation for classical Be stars both have a common origin. The excess is due to free-free and free-bound emission from the dense ionized CS gas, with the polarization being engendered by scattering within clouds of free electrons around the star. Moreover, classical Be stars are very fast rotators and at least a few of them rotate close to their break-up velocities. This rapid rotation is considered as a trigger mechanism for CS envelope formation. To obtain some conclusions on the geometry of CS envelopes around classical Be stars, analysis of correlations between different observational parameters seems strongly desirable (see, for example, Coté & Waters 1987). For the same reason, it is very important to investigate each of the observed parameters statistically for as large a representative sample as possible. During the last 20 years there have been several attempts to carry out such kinds of investigation for large samples of classical Be stars (from 50 up to 200 objects). However, such sample sizes are still sometimes insufficiently large to provide irrefutable conclusions and some authors indicate this fact themselves (see, for example, Quirrenbach et al. 1997 or Ghosh et al. 1999). By taking a relatively large sample of 90 objects, by adding the "new'' data from Ghosh et al. (1999) and McDavid (1999, 2000) to the "old'' data from McLean & Brown (1978) and Poeckert et al. (1979), Yudin (2000) came to quite different conclusions on the dependence between $\,p_{\ast}\,$and $\,v\sin{i}\,$ than those of McLean & Brown (1978) who used only 67 stars.

The general aim of the proposed exercises here is to investigate statistically the data of $\,p_{\ast}\,$ and $\,v\sin{i}\,$ for the largest possible sample of classical Be stars, to compare the behaviour with near and far IR excesses and to search for possible correlations between the parameters. Another aim is to investigate possible differences between spectral subgroups from the general assembly of classical Be stars. These analyses may give new insight on the configuration of CS envelopes.

2 Selection procedure

Up till now, there is no statistical information on whether or not classical Be stars can be treated as a homogeneous group. The largest database, SIMBAD, contains only 82 objects which are classified as Be stars but, at the same time, 884 objects are classified as O-B emission line stars. The largest revised catalogue of Be stars of Jaschek & Egret (1982) contains 1155 objects but a significant fraction of these belong to the young Herbig Be star group.

For the statistical analysis here, data for 627 Be stars have been compiled from different sources, taking into account an overlap of the list of emission line stars from the SIMBAD database and the catalogue of Be stars of Jaschek & Egret (1982). All Be stars which have been considered in the past as young Herbig Be stars have been excluded; a few Be stars from recent papers (see, for example, Steele et al. 1999) have been added. Of course it is impossible to be sure that all the collected objects are definitely of classical type, but the fraction of "doubtful" objects cannot be large (as follows from further discussion, see Sect. 6 and Fig. 2). Thus we can say that the compilation represents the largest ever, more or less homogeneous, sample of classical Be stars. Note that for selection, we collected only stars of III, IV and V luminosity class and the sample does not include any supergiants. Emission line stars of the spectral type A and later are also excluded. No limitation has been made on the level of observed polarization ( $p_{\rm obs}$) of the sample stars. Thus we have in hand for statistical study:

1.

    497 stars with polarization data;

2.

    246 stars with values of near IR excess E(V-L);

3.

    151 stars with values of IR excess at 12 $\mu$m, and

4.

    463 stars with $\,v\sin{i}$ values.

All collected data are presented in Appendix I. It is, of course, of direct interest to investigate the distributions of the observed parameters for the entire sample, but in addition, it is also important to compare the data according to different spectral subgroups. The latter is obviously important to astrophysical understanding but such breakdown also allows investigation of the homogeneity of the total assembly. To obtain statistically representative data sets, we separate the entire sample into four subgroups. Subdividing our list yields 174 stars in the interval O-B1.5, 170 stars in the interval B2-B2.5, 147 stars in the interval B3-B5.5, and 135 stars in the interval B6-B9.5. This separation was not justified physically a priori, but it reflects an approximately equal range of effective temperature $T_{\rm e}$( $\Delta T_{\rm e}\approx5000$ K) for the last three subgroups.

3 Data reduction

Because the main aim of the present investigation is a statistical study of different observed parameters and correlations between them, it is important to minimize any uncertainties and all kind of statistical biases. Since a principal characteristic of Be stars is their temporal variability, it is best to use the mean values of investigated parameters for each star. The averaging was done by the weighting procedure according to the standard errors of individual measurements (see, for example, Smart 1958).

For the discussion of the various parameters, it is useful to construct the relative distributions of observed values in the form of histograms. For this reason it is very important to choose the best binwidths for the distributions. As will be shown below, the chosen binwidths of 0.3%, $0\mbox{$\,.\!\!\!^{\rm m}$ }3$ and 50 kms^-1 for polarization, colour and rotational velocity respectively are typically larger than the error associated with the individual measurements and even larger than the range of variability associated with any of the stars. Consider now the reduction of each of the recorded parameters.

3.1 Reduction of v sin i

At present, most Be star investigations have used the data of $\,v\sin{i}\,$ mainly from Slettebak (1982) (164 classical Be stars). Other old data are also available and these are supplemented by numerous more recent collections. The only problem in compiling $\,v\sin{i}\,$ values from different sources is that they have different scales and accurate calibrations between the compiled data are required to avoid systematic errors. It is well known that there is a difference between the "new'' and "old'' Slettebak's scales (Slettebak et al. 1975). First of all we compared the $v\sin{i}$ data of Slettebak (1975) with those determined or compiled by others for the same stars and calculated least-squares fits to the data. Following this calibration process, we transformed individual values of $\,v\sin{i}\,$ from different sources into the common "new'' scale of Slettebak et al. (1975). Stars which after transformation provided strong disparities in their values according to the source catalogue were excluded. Average weighted values of $\,v\sin{i}\,$ and associated standard errors were obtained for each star using the weighting procedure already mentioned above (see Smart 1958). For the case of single measurements, we used the value of a standard error indicated in the reference or adopted a value for the error equal to 10% of the measured value. This procedure allows exclusion of most systematic and incidental errors and helped to obtain a new expanded homogeneous data set. The determined values of $\,v\sin{i}\,$ for each star with the standard error are presented in Appendix I. It can be seen that for more than 95% program stars, the standard error is less than $\pm 40$ kms^-1.

3.2 Reduction of $\mathsfsl{E(V-L)}$

The values of near IR excesses E(V-L) were calculated by a classical approach using the measurements of observed and normal colour indices and interstellar reddening in the respective photometric bands, this being summarized as:

$$E(V-L)=(V-L)_{\rm obs}-(V-L)_0-(A_{V}-A_{L})$$ (1)

where $(V-L)_{\rm obs}$ and (V-L)₀ are the observed and normal colour indices and $\,A_{V}$, $A_{L}\,$ are the values of interstellar reddening in the respective photometric bands. We assume that no circumstellar dust environment exists around most of the considered stars and E(V-L) is the excess due to the contribution of circumstellar gas. For most stars the values of $\,A_{V}$, $A_{L}\,$ were calculated assuming a normal extinction law in the form $\,A_{V}=3.1\times E(B-V)\,$, $\,A_{L}\approx0.06A_{V}\,$ and $\,E(B-V)=(B-V)_{\rm obs}-(B-V)_{0}\,$(for a detailed discussion, see, for example, Dougherty et al. 1994). In principle, this procedure may lead to the over-dereddening but the error in most cases does not exceed $\,0\mbox{$\,.\!\!\!^{\rm m}$ }2\,$and is unimportant, taking into account the chosen histogram binwidth for E(V-L) of $0\mbox{$\,.\!\!\!^{\rm m}$ }3$. The values of (V-L)₀ and (B-V)₀ were taken from Straizys (1977). It is easy to show that due to the uncertainty $\sigma_{E(B-V)}\approx0\mbox{$\,.\!\!\!^{\rm m}$ }05$, the uncertainty in E(V-L) is about $0\mbox{$\,.\!\!\!^{\rm m}$ }2$. Of greater importance is the study of errors due to the range of photometric variability for individual objects. Most classical Be stars are photometric variables on different time scales. Hubert & Floquet (1998) recently investigated photometric variability of a sample of 289 classical Be stars using the Hipparcos photometric data. They found that most of their program stars (about 33%) exhibited photometric variability $0\mbox{$\,.\!\!\!^{\rm m}$ }1\leq\Delta H_{\rm p}\leq0\mbox{$\,.\!\!\!^{\rm m}$ }3$. Other stars usually show a lower level of photometric variability. Similar conclusions were made recently by Moujtahid et al. (1998) who compiled the values of $\,\sigma_{V}\,$for 50 well known classical Be stars. Moreover, for the calculation of E(V-L), optical and IR data were selected to be as close as possible in time.

3.3 Reduction of $\mathsfsl p_{\ast}$

Currently, the intrinsic polarization components have been estimated for only about 90 classical Be stars.

	Figure 1: Dependence between the interstellar polarization and the distance in the $5\hbox {$^\circ $ }$ field surrounding a few program stars. The numerical values attached to the curves refer to the position angle of interstellar polarization according to distance
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In order to obtain a $p_{\ast }$ value for each star, its interstellar component, $p_{\rm is}$, needs to be subtracted. This is best done using the "field star method" whereby the local growth of $p_{\rm is}(D)$ with distance, (D), is first obtained from measurements of stars close to the line of sight of the target star. According to the distance of the considered Be star, an estimate for its $p_{\rm is}$ component is then made from the calibration and vectorially subtracted using the component Stokes parameters. The selected fields ranged from $1\hbox{$^\circ$ }$ to $5\hbox {$^\circ $ }$ depending on the number of measured stars in the vicinity. A recent advance for the method has been the improvement in accuracy by the application of Hipparcos data (at least for objects with D<500 pc), this being used in the reduction scheme for most of the classical Be stars in our list. Fortunately, the average weighted values of observed polarization, $p_{\rm obs}$, for most objects in our list together with the associated $\,\sigma_{p_{\rm obs}}$were calculated recently by Heiles (2000), so considerably reducing the amount of laborious work required for the exercise. For a few stars which were investigated polarimetrically more intensively, we estimated $\sigma_{p_{\rm obs}}$, taking into account the data from recent papers of Quirrenbach et al. (1997), McDavid (1999, 2000) and Ghosh et al. (1999).

Two typical examples of the reduction scheme are presented in Fig. 1. Note that, in some cases, $\,p(D)\,$ may not display monotonic growth due to the line of sight passing through interstellar dust clouds with complex structure and different orientation of dust grains. For the cases where the distance to an object could not be determined from the Hipparcos data or the parallax errors were too large, we estimated the distance to a star either from reddening or from an absolute magnitude-spectral type/luminosity calibration. In a few cases an interstellar component has been estimated by the investigation of the uniformity of polarization vectors in the vicinity of the star. The parameters of $p_{\rm obs}$, $p_{\rm is}$ and $p_{\ast }$ for the program stars are presented in Appendix II. It is seen that except for a few stars, most values of $\sigma_{p_{\rm obs}}< 0.3\%$ and thus less than the chosen histogram binwidth.

An independent by-product of this study is the new set of estimates of interstellar polarization components, this being useful for other future investigations.

4 $\mathsf{E(V-L)}$, polarization and v sin i as a function of spectral type

4.1 The excess $\mathsfsl{E(V-L)}$ against spectral type

In Fig. 2a the values of the excess E(V-L) are plotted against spectral type for the entire sample. The following features are immediately obvious. Firstly, there is a clear upper limit for a given spectral type with the maximum value of about $1\mbox{$\,.\!\!\!^{\rm m}$ }8$ for B0-B2 stars and with the decrease of near IR excesses toward later than B2 spectral types. A similar result was obtained by Dougherty et al. (1991) who found that most of late type stars show little or no near IR excess colour but for a significantly smaller statistical sample. However, in contrast to previous studies, we conclude that the decrease of the upper limit of E(V-L)towards late spectral types is well fitted by a function in the form of an exponential decay rather than the linear decay. For the spectral interval from B1 to B9.5 the best fit is described as follows:

$$E(V-L)=0.31\,+\,1.67\times\exp\left(-\frac{X-0.70}{2.94}\right)$$ (2)

where $\,X\,$ is the value of the spectral subtype. This fit is indicated by a solid curve in Fig. 2a with 95% of program stars of a given spectral types located below the curve. Note that approximately the same curve represents the change of the stars' effective temperatures according to spectral subtype. This behaviour is readily understandable, in general terms, as the near IR excess is dependent on the electron density in the circumstellar environment which is in turn a function of stellar temperature. Secondly, the value of IR excesses shows a larger scatter for B1-B3 spectral types. Thirdly, less than 5% of the investigated stars are above the limit (2). The discussion on the anomalous objects and the distributions of IR excesses for different spectral subgroups of Be stars will be considered in more detail in Sect. 6.

	Figure 2: Near IR excess E(V-L), intrinsic polarization, and $v\sin{i}$as a function of spectral type. The upper limit is indicated by the solid curve. The short-dashed curve indicates the average value of the parameters for a given spectral type. The long-dashed curve indicates the value of $T_{\rm e}$ for a given spectral type
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4.2 Polarization against spectral type

In Fig. 2b the values of $\,p_{\ast}\,$ for the entire sample are plotted against spectral type. This distribution shows a similar behaviour as for near-IR excesses. Firstly, the values show larger scatter for Be stars of earlier spectral types (B0-B3). Secondly, there is an upper limit of $\,p_{\ast}\,$ for a given spectral type. This is represented by a solid curve in Fig. 2b below which 95% of stars of the given spectral type (in the range B3-B9) are located. The upper limit is well described by a function in the form of an exponential decay:

$$p_{\ast}(\%)=0.46\,+\,1.74\times\exp\left(-\frac{X-1.5}{3.22}\right)$$ (3)

where $\,X\,$ is the number defining the spectral subtype. The explanation of this dependence is the same as in the previous section. Thirdly, on average, Be stars of earlier spectral type exhibit larger values of polarization. Fourthly, the maximum values of intrinsic polarization for B0-B3 stars are approximately the same ( $p_{\ast}\approx2$%) and the upper limit in polarization for these stars does not correspond to the curve which represents the change of star's effective temperature according to spectral subtypes. This indicates that the behaviour "$p_{\ast }$: Sp class" for early type stars is more complicated and there are several physical parameters dictating the levels of polarization. The distributions of $\,p_{\ast}\,$ for different spectral subgroups of Be stars will be considered later in Sect. 7.

4.3 v sin i against spectral type

In Fig. 2c the values of $\,v\sin{i}\,$ are plotted against spectral type for the entire sample. First impressions suggest that there are no significant differences in $\,v\sin{i}\,$ for the subtypes of Be stars. However, there is a possible trend in the average values in the sense that later spectral type stars rotate faster than early ones. The best linear fit for this trend is:

$$v\sin{i}=(204\pm5)+(2.2\pm1.2)X$$ (4)

where $\,X\,$ again is the value of the spectral subtype. Note, however, that this trend is not obvious for stars of III luminosity class and it is more pronounced for IV and V classes:

$${\rm III}\!:~v\sin{i}=(190\pm14)-(1.1\pm2.8)X$$

$${\rm IV}\!\!:~v\sin{i}=(183\pm14)+(8.0\pm3.5)X$$

$$\ {\rm V}\!\!:~v\sin{i}=(205\pm~7)+(3.8\pm1.6)X.$$

The correlations for the above dependencies are small or even absent (correlation coefficients: $r=0.09\pm0.06$ for the entire sample, $-0.05\pm0.12$, $0.27\pm0.11$ and $0.14\pm0.06$ for the III, IV and V classes respectively).

The value of projected rotational velocity shows a large scatter for all Be star subspectral types but with a clear concentration around the values between 200 to 250 kms^-1. All Be stars exhibit a projected rotational velocities which are less than a critical rotational velocity for the given spectral type. We will investigate the distribution of $\,v\sin{i}\,$ in more detail in the next section.

5 Distributions of projected rotational velocities

Two well-known investigations of projected rotational velocities of Be stars were completed 18 years ago (Slettebak 1982 - for a sample of 164 objects; Fukuda 1982 - for a sample of 239 objects). Slettebak noted that there are no strong differences in the rotational characteristics of Be stars of different spectral types or luminosity classes. To the contrary, Fukuda arrived at a different conclusion showing that the middle and later types rotate faster on average than early Be stars, but the influence of luminosity class was not studied. Since then, several attempts to carry out similar kinds of investigation for different samples of classical Be stars have been made with somewhat different conclusions (see, for example, Zorec et al. 1990; Balona 1990; Zorec & Briot 1997; Grady et al. 1989; Steele 1999 etc.). It is now possible to reassess the situation using the numerous and homogeneous $\,v\sin{i}\,$ data collected and reduced in the present study. We investigate here the distributions of projected rotational velocities of the program stars for four different spectral subgroups (O-B1.5, B2-B2.5, B3-B5.5 and B6-B9.5) and different luminosity classes.

Investigation of a distribution such as $\,v\sin{i}\,$ is a classical problem of mathematical statistics. In the first place, we corrected the constructed histograms from the influence of errors ( $\sigma_{v\sin{i}}$) using Eddington's algorithm (see Smart 1958). One of the distributions (viz: for the entire sample of Be stars) is presented in Fig. 3. To compare the distributions, the following parameters should be calculated and analyzed: the mean ( $\overline{v\sin{i}}$) and the root mean square deviation ($\sigma$). To calculate these parameters for a given distribution we used the weighting procedure according to the standard error associated with each contributing measurement (see, for example, Brooks et al. 1994). However, in order to make meaningful comments on the similarity or differences between the distributions, in the first place their distribution functions should be analyzed. As a first step we should try to investigate their "Normality''. If the distribution is demonstrably non-Normal, other types of function might then be considered and applied. To test for Normality, we calculated the coefficients of skewness (g₁) and the coefficients of kurtosis (g₂) for each of the distributions. A detailed description of this procedure has been given recently by Brooks et al. (1994) together with the analytical expressions for $\,g_{1}\,$ and $\,g_{2}\,$. We calculated unbiased estimators of these coefficients with the associated variances (see Table 1) and compared the values with those tabulated in the above mentioned paper for the 95% and 99% confidence levels. This analysis indicates immediately that all of the distributions may be considered as Normal at the 99% confidence level and thus we don't need to invoke other types of distribution functions. At this stage, we cannot say whether the high arithmetical significance of the statistics for Normality carries real physical significance in the sense that it is an expected intrinsic characteristic of the $v\sin{i}$ distributions of Be stars. Such kind investigations might be the matter of future studies.

	Figure 3: Observed distribution of $v\sin{i}$ corrected from errors and its best Gaussian fit for the entire sample of Be stars. The dashed vertical line indicates the mean value of $v\sin{i}$
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The calculated parameters for each of the distributions are given in Table 1.

 

 

Table 1: Mean $\overline{v\sin{i}}$ values, $(\sigma)$, the coefficients of skewness (g₁) and kurtosis (g₂)of the distributions, true rotational velocities and the ratio of $\overline{v\sin{i}}/\overline{v_{\rm c}}$ and $\overline{v_{\rm t}}/\overline{v_{\rm c}}$ for different spectral subgroups and different classes of luminosity of Be stars

III, III-IV classes luminosity
spectral range:	O-B1.5	B2-B2.5	B3-B5.5	B6-B9.5	O-B9.5
N stars:	37	30	28	40	135
$\overline{v\sin{i}}$, (kms-1)	$203\pm11$	$211\pm10$	$192\pm14$	$207\pm10$	$201\pm~6$
$\sigma$	$70\pm~5$	$78\pm~4$	$89\pm~6$	$63\pm~5$	$74\pm~8$
g1	$-0.326\pm0.158$	$0.314\pm0.162$	$0.137\pm0.182$	$-0.234\pm0.143$	$-0.001\pm0.042$
g2	$1.998\pm0.605$	$2.441\pm0.621$	$2.454\pm0.693$	$2.653\pm0.549$	$2.412\pm0.164$
$\overline{v_{\rm c}}$, (kms-1)	456	408	383	320	385
$\overline{v\sin{i}}/\overline{v_{\rm c}}$	0.45	0.52	0.50	0.65	0.52
$\overline{v_{\rm t}}$, (kms-1)	259	269	245	264	259
$\overline{v_{\rm t}}/\overline{v_{\rm c}}$	0.57	0.66	0.64	0.83	0.67
IV-V, V classes luminosity
spectral range:	O-B1.5	B2-B2.5	B3-B5.5	B6-B9.5	O-B9.5
N stars:	66	78	78	61	283
$\overline{v\sin{i}}$, (kms-1)	$202\pm~9$	$207\pm~8$	$236\pm~7$	$243\pm~8$	$229\pm~4$
$\sigma$	$80\pm~9$	$77\pm~6$	$41\pm~6$	$58\pm~5$	$72\pm~7$
g1	$0.188\pm0.087$	$0.040\pm0.074$	$-0.586\pm0.074$	$-0.353\pm0.094$	$-0.195\pm0.021$
g2	$2.924\pm0.339$	$2.808\pm0.290$	$4.005\pm0.290$	$2.802\pm0.365$	$2.863\pm0.083$
$\overline{v_{\rm c}}$, (kms1)	527	479	448	398	464
$\overline{v\sin{i}}/\overline{v_{\rm c}}$	0.38	0.43	0.53	0.61	0.49
$\overline{v_{\rm t}}$, (kms-1)	257	264	300	309	292
$\overline{v_{\rm t}}/\overline{v_{\rm c}}$	0.49	0.55	0.67	0.78	0.63
all classes luminosity (CL) including the stars with undefenite CL
spectral range:	O-B1.5	B2-B2.5	B3-B5.5	B6-B9.5	O-B9.5
N stars:	106	116	121	120	463
$\overline{v\sin{i}}$, (kms-1)	$201\pm~7$	$205\pm~7$	$241\pm~6$	$225\pm~6$	$219\pm~3$
$\sigma$	$74\pm~4$	$72\pm~5$	$59\pm~5$	$57\pm~4$	$70\pm5$
g1	$0.392\pm0.055$	$0.110\pm0.050$	$-0.400\pm0.048$	$-0.234\pm0.049$	$-0.043\pm0.013$
g2	$3.333\pm0.216$	$2.666\pm0.198$	$2.995\pm0.191$	$2.971\pm0.192$	$2.877\pm0.051$
$\overline{v_{\rm c}}$, (kms-1)	502	458	429	362	436
$\overline{v\sin{i}}/\overline{v_{\rm c}}$	0.40	0.45	0.56	0.62	0.50
$\overline{v_{\rm t}}$, (kms-1)	256	261	307	288	280
$\overline{v_{\rm t}}/\overline{v_{\rm c}}$	0.51	0.57	0.72	0.80	0.64
the data from Slettebak 1982
spectral range:	B0-B1.5	B2-B2.5	B3-B5.5	B6-B9.5
N stars:	25	35	51	54
$\overline{v\sin{i}}$, (kms-1)	236	193	221	228
the transformed data from Fukuda 1982
spectral range:	B0-B1.5	B2-B2.5	B3-B5.5	B6-B9.5	O-B9.5
N stars:	69	63	50	57	239
$\overline{v\sin{i}}$, (kms-1)	189	196	239	244	214
the data from Zorec & Briot 1997
spectral range:	B0-B1.5	B2-B3	B4-B7	B8-B9.5
$\overline{v_{\rm t}}$, (kms-1) (II-III, III, III-IV classes)	230	190	220	130
$\overline{v_{\rm t}}$, (kms-1) (V class)	240	250	250	280
the data from Grady et al. 1989
spectral range:				B6-B9.5
N stars:				40
$\overline{v\sin{i}}$, (kms-1)				226

 

 

Table 2: The values of $v_{\rm c}$ for Be stars of different spectral subtypes and class luminosity

Spectral type:	O8	O9	B0	B0.5	B1	B1.5	B2	B2.5	B3	B4	B5	B6	B7	B8	B9
$v_{\rm c}$, III class (kms-1):	500	460	430	412	395	383	365	355	340	325	310	295	284	275	270
$v_{\rm c}$, IV class (kms-1):	600	560	520	502	483	465	445	435	420	400	386	374	362	352	344
$v_{\rm c}$, V class (kms-1):	626	596	565	540	520	500	483	470	460	442	425	410	400	396	391

It is of interest also to compare the parameters of distributions of true rotational velocities which are easily calculated by the method suggested by Chandrasekhar & Münch (1950). In the case of randomly oriented stellar rotation axes, the mean true rotational velocity may be calculated by a simple equation: $\overline{v}_{\rm t}=\frac{4}{\pi}\overline{v\sin{i}}$.

To test the differences between the dispersions, we can use Fisher's criterion (F-test). Simple calculations indicate that for the III-IV luminosity classes, the differences between the dispersion for the distributions are statistically insignificant at the 95% confidence level (except late type stars). For the V class, it was found that the difference in $\sigma$ is significant at 95% confidence level between O-B2.5 and B3-B9 spectral subgroups.

As the next step, we may compare the mean values of $\overline{v_{\rm t}}$ and the ratio of $\,\overline{v_{\rm t}}/v_{\rm c}\,$ for selected subgroups. It is well known that the mean true velocities and critical velocities are very different, depending on the luminosity class and spectral class (see for example Zorec et al. 1990). For our purposes we interpolated/extrapolated the values of $\,v_{\rm c}\,$ from Moujtahid et al. (1999) and Zorec (2000) and present them in Table 2. The values of $\overline{v_{\rm c}}$ for given spectral subgroups (presented in Table 1) have been calculated taking into account the specific population of stars of different spectral subtypes and luminosity classes in a subgroup. Simple calculations indicate that the differences between the mean values of $\overline{v_{\rm t}}$ (or $\overline{v\sin{i}}$) for all spectral subgroups of III-IV luminosity classes are insignificant. There is, however, a pronounced trend in the ratio $\,\overline{v_{\rm t}}/v_{\rm c}\,$ for these stars. For the V class, no differences were found between the O-B1.5 and B2-B2.5 subgroups, as well as between subgroups B3-B5.5 and B6-B9.5. Nevertheless, for the subgroups O-B2.5 and B3-B9.5, there is strong statistically significant difference at the 99% confidence level in agreement with the result of Fukuda (1982) who found a clear distinction in $v\sin{i}$between B2 and B3 stars. For reference, we present in Table 1 the $\overline{v\sin{i}}$ data from other sources. The data of Fukuda (1982) (from his Table 5) were transformed into the "new'' Slettebak's scale. As follows from Table 1, most early determined values are in agreement (in general) with our data (obtained with 2 to 3 times better statistical quality).

We confirm that late type Be stars rotate faster on average than early types and the trend in the ratio $\,\overline{v_{\rm t}}/v_{\rm c}\,$ (see Table 1) is also real. We conclude that late type Be stars of all luminosity classes rotate closer to their critical break-up velocity. Whether these differences have affect on the physical characteristics of CS envelopes, or not, will be considered in Sect. 8.

Finally, it was found that the ratio $\overline{v\sin{i}}/v_{\rm c}$ (or $\,\overline{v_{\rm t}}/v_{\rm c}\,$) increases from dwarfs to sub-giants/giants for early spectral type stars (O-B1.5 and B2) though not obvious for middle and late spectral types (see Table 1). This clearly contradicts the results of Steele (1999) who investigated a small sample of Be stars separated into three subgroups according to the luminosity class (III, IV, V-13, 8 and 34 objects respectively). Although our statistical study for sub-giants and giants is also not based on a large sample (135 objects), the behaviour is clearer for the sequence of ${\rm V}\rightarrow {\rm IV}\rightarrow {\rm III}$ luminosity classes (see Table 3). This phenomenon can be easily explained in terms of a more rapid evolution of high-mass stars and the acceleration of their rotation. The main sequence lifetime is significantly shorter for B0 stars than for B9 stars (see, for example, Bisnovatyi-Kogan 1989). It may be supposed that most of O-B1.5 and B2 stars of III and IV luminosity classes are possibly in the secondary contraction phase but, on average, the late spectral type stars are not.

 

 

Table 3: The values of $v\sin{i}/v_{\rm c}$ for Be stars of different spectral subtypes and classes luminosity

Spectral type:	O-B1.5	B2-B2.5	B3-B5	B6-B9
$v\sin{i}/v_{\rm c}$, V class:	0.38	0.43	0.53	0.61
$v\sin{i}/v_{\rm c}$, IV class:	0.43	0.49	0.50	0.64
$v\sin{i}/v_{\rm c}$, III class:	0.46	0.56	0.51	0.65

6 Distributions of near and far IR excesses

As already mentioned, a detailed study of near IR excesses of classical Be stars was undertaken by Dougherty et al. (1994). They used near IR excesses for a sample of 144 Be stars (115 stars with E(V-L) data) for statistical analysis. One of the conclusions was that the excess increases with wavelength and 60% of Be stars exhibit a significant excess in the L band. They also found a bi-modal distribution of the excesses for the wavelengths greater than 2.2 $\mu$m but they did not study this behaviour in detail. For this reason, we investigate here the distributions of E(V-L)colour excesses. We have collected the data of V and L photometry for $\approx$200 stars and, for an additional 40 stars, the values of E(V-L)were estimated by an extrapolation of JHK data. Thus, by essentially doubling the size of the data base over Dougherty et al. (1994), we are able to compare the E(V-L) distributions for different subgroups of Be stars.

	Figure 4: Observed distributions of E(V-L) (corrected for errors) for different subgroups of Be stars. The short-dashed and long-dashed vertical lines indicate the positions of two peaks for a bi-modal Gaussian distribution
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Note, however, that 6 stars were excluded from the statistical analysis due to their unexpectedly large near IR excess for their spectral type (see Fig. 2a). Some comments on their characteristics and evolutionary status are given below. First of all, the values of near IR excess, $E(V-L)>2^{\rm m}$, are difficult to explain in a framework of free-free, free-bound emission and such large values require the presence of a dust component in CS environment. Thus, the star HD 181615 (B2Vpe, $E(V-L)=3\mbox{$\,.\!\!\!^{\rm m}$ }03$) is likely to contain CS dust. The conclusion on the presence of the dust emission is very likely for the stars HD 50123 (B6IVe, $1\mbox{$\,.\!\!\!^{\rm m}$ }42$), HD 51480 (B8e, $1\mbox{$\,.\!\!\!^{\rm m}$ }53$), both exhibiting also an excess at 12 $\mu$m larger than the upper limit for Be stars of a given spectral type (Coté & Waters 1987). Note that 51 Oph (B9.5V), HD 164906 (B1IV:) and HD 38087 (B5Ve), exhibit unusual far IR excess ( $3\mbox{$\,.\!\!\!^{\rm m}$ }7$, $4\mbox{$\,.\!\!\!^{\rm m}$ }98$ and $3\mbox{$\,.\!\!\!^{\rm m}$ }79$ at 12 $\mu$m respectively). Relative to classical Be stars, this also indicates the presence of dust emission. An inspection of the lists of Coté & Waters (1987) and Oudmaijer et al. (1992) shows that only a few Be stars from about 150 objects show excess emission with $E(V-12)>3^{\rm m}$ and some of them are definitely young objects. Thus, the value $E(V-12)\approx3^{\rm m}$ is a limit for free-free and free-bound emission in the CS environment of classical Be stars. The evolutionary status of two other objects: HD 177291 (B8Ve, $1\mbox{$\,.\!\!\!^{\rm m}$ }44$) and HD 81357 (B8Ve, $1\mbox{$\,.\!\!\!^{\rm m}$ }5$) is not clear. However, the latter star has been classified recently as a binary with an F/G secondary (Halbedel 1996) whose light contribution may give an increase to the observed E(V-L) excess.

The distribution of E(V-L) for the entire sample shows the same bi-modal shape behaviour as described by Dougherty et al. (1994) but in a more pronounced form. We performed the same procedure for error correction of the histograms as in the previous section.

 

 

Table 4: Parameters of distributions of near IR E(V-L) and far IR E(V-12) excesses (in magnitudes) for different subgroups of Be stars

Subgroup:	O-B1	B2	B3-B5	B6-B9
N stars:	57	59	58	67
$\overline{E(V-L)}$	0.93	0.72	0.41	0.17
N stars:	29	38	38	40
$\overline{E(V-12)}$	1.90	1.59	1.39	0.72
Parameters of bi-modal Gaussian fits for E(V-L)
Peak 1: $\bar{x}_{1}$	$0.32\pm0.05$	$0.24\pm0.03$	$0.16\pm0.02$	$0.16\pm0.01$
Peak 1: $\sigma_{1}$	$0.12\pm0.02$	$0.14\pm0.03$	$0.20\pm0.01$	$0.16\pm0.02$
Peak 2: $\bar{x}_{2}$	$1.30\pm0.07$	$0.92\pm0.04$	$0.80\pm0.05$
Peak 2: $\sigma_{2}$	$0.37\pm0.08$	$0.37\pm0.05$	$0.11\pm0.04$
Parameters of bi-modal Gaussian fits for E(V-12)
Peak 1: $\bar{x}_{1}$	0.75 ?	$0.80\pm0.05$	?	$0.71\pm0.06$
Peak 1: $\sigma_{1}$		$0.27\pm0.04$		$0.30\pm0.04$
Peak 2: $\bar{x}_{2}$	$2.07\pm0.06$	$2.05\pm0.04$	$1.83\pm0.05$
Peak 2: $\sigma_{2}$	$0.27\pm0.05$	$0.32\pm0.04$	$0.30\pm0.07$

Firstly, late spectral type stars exhibit smaller values of E(V-L) on average (see Table 4), this already mentioned in Sect. 4 and noted earlier by Dougherty et al. (1991, 1994), possibly due to a lower electron density in their CS environment. It is obvious from Fig. 4 and in Table 4, that a bi-modal form is present for the first three spectral subgroups of Be stars (O-B1.5, B2-B2.5 and B3-B5.5). For these subgroups, the mean values of E(V-L) corresponding to the first peak are approximately the same, within the length of the sampling step, with $\overline{E(V-L)}_{\rm peak 1}\approx0\mbox{$\,.\!\!\!^{\rm m}$ }22\pm0\mbox{$\,.\!\!\!^{\rm m}$ }08$. Statistical analyses, similar to those made in previous section, indicate that all the distributions may be well approximated by a bi-modal Gaussian form. This bi-modality for O-B1.5, B2-B2.5 and B3-B5.5 spectral subgroups is confirmed to a 99% level of confidence. We performed a more detailed study of the B3-B5.5 spectral subgroup and found that a bi-modal distribution is present for B3 type stars with two peaks centered at $0\mbox{$\,.\!\!\!^{\rm m}$ }15\pm0\mbox{$\,.\!\!\!^{\rm m}$ }16$ and $0\mbox{$\,.\!\!\!^{\rm m}$ }87\pm0\mbox{$\,.\!\!\!^{\rm m}$ }21$), but bi-modality is not present for B4-B5 stars. For B6-B9.5 stars, the distribution also does not exhibit bi-modality and the mean value of E(V-L) for the best simple Gaussian fit is the same as the first peaks for other subgroups.

	Figure 5: Observed distributions of E(V-12) (corrected from errors) for the entire sample of Be stars. The short-dashed and long-dashed lines indicate the position of two peaks for a bi-modal Gaussian distribution
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	Figure 6: Observed distributions of E(V-12) (corrected from errors) for different subgroups of Be stars. The short-dashed and long-dashed lines indicate the positions of the two peaks for a bi-modal Gaussian distribution
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The important difference between the various bi-modal distributions is that the second peak has its maximum at different values of E(V-L). The behaviour is that for early spectral types, the second peak appears at larger values of E(V-L). This difference is clearly obvious in a numerical form in Table 4 where the parameters of these distributions are given. Using a t-test and T-test, it is easy to demonstrate that the differences between the positions of the second peaks are statistically significant at the 99% confidence level. Note that, for the optically thin case, the excess emission in the near IR is proportional to the emission measure which is a function of the mean electron density in formation regions and the ionizing stellar fluxes. It is easy to show from our data that the changes in $E(V-L)_{\rm peak 2}$ may be well described by a linear function of $\log\overline{T_{\rm eff}}$ in the form: $E(V-L)_{\rm peak 2}\approx{-11.1}+2.78\log\overline{T_{\rm eff}}$. The predicted $E(V-L)_{\rm peak 2}$ values correspond well to the observed values.

	Figure 7: E(V-12) colour excess versus spectral type for the entire sample of Be stars. The upper limit is indicated by the solid line. The long-dashed line indicates the value of $T_{\rm e}$ for a given spectral type
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It is of interest to compare the behavior of E(V-L) with that for far IR excesses at $12~\mu$m. Dougherty et al. (1994), with the reference to a private communication of Waters, mentioned the presence of bi-modality of the E(V-12) distribution for Be stars. To check this suggestion we constructed the similar diagrams for E(V-12). Although bi-modality is clearly seen for the entire sample of Be stars (Fig. 5), it reflects rather the different contributions of early and late type stars. Our sample is insufficiently large to provide any conclusion on bi-modality for different spectral subgroups. Slight evidence of two peaks in the E(V-12) distribution exists only for B2 stars (see Fig. 6). An important point, however, is that the average value of $\overline{E(V-12)}$ and the position of the maximum peak in the distributions is not significantly different between the first three subgroups (see Table 4). Coté & Waters (1987) considered the far IR data for a sample of 101 Be stars and found the constant upper limit in E(V-12) for stars earlier than B3. The same conclusion follows from the analysis of our data for 150 objects (see Fig. 7). This indicates that the characteristics of the outer parts of CS envelopes are similar for early and middle spectral type stars. For the majority of Be stars, with the excess emission increasing with wavelength, the CS envelopes must become optically thick from at least $\lambda > 10~\mu$m (Gehrz et al. 1974). With this assumption, for any kind of geometry for the CS environment, the excess E(V-12) is proportional to the effective surface area of the emitting region. It is easy to show that the relation between the upper limit of the excess E(V-12)and $T_{\rm e}$ will be fairly linear for B3-B9 stars in contrast to the E(V-L): $T_{\rm e}$ relation. But the behaviour in early type stars is more complicated.

An unresolved point is the nature of the peak 1 in the E(V-L) distributions for early and intermediate spectral type stars. There are several possibilities to explain the value of $E(V-L)\approx 0$ for some objects:

a) there are many Be stars surrounded by small CS envelopes indicated by low emission in H$\alpha$;
b) the photometric data of V and L used for calculation of the E(V-L) excesses were obtained mainly non-simultaneously;
c) spectral classification of some early type stars can be dubious.

For case (a), we would expect a lower $\,p_{\ast}\,$ values for the stars which are concentrated under peak 1. Although the analysis of our data shows some decrease in the polarization level for such stars in comparison with peak 2, the difference in $\overline{p_{\ast}}$ for the two peaks is small and insignificant (see Table 5).

 

 

Table 5: Differences in the level of intrinsic polarization for the stars in the peak 1 and peak 2 of the E(V-L) distributions

Subgroup:	O-B1		B2		B3-B5
	peak 1	peak 2	peak 1	peak 2	peak 1	peak 2
	$E(V-L)<0\mbox{$\,.\!\!\!^{\rm m}$ }85$	$E(V-L)>0\mbox{$\,.\!\!\!^{\rm m}$ }85$	$E(V-L)<0\mbox{$\,.\!\!\!^{\rm m}$ }4$	$E(V-L)>0\mbox{$\,.\!\!\!^{\rm m}$ }5$	$E(V-L)<0\mbox{$\,.\!\!\!^{\rm m}$ }4$	$E(V-L)>0\mbox{$\,.\!\!\!^{\rm m}$ }5$
N stars:	26	26	16	38	26	20
$\overline{p_{\ast}}$, %	0.57	0.85	0.52	0.63	0.42	0.52

For the cases (b) and (c), we would expect a more single-peaked distribution with more extended tails rather than a bi-modal distribution. Finally, it should be remembered that each Be star can show at any time one of the phases - normal B $\Leftrightarrow$ Be $\Leftrightarrow$ Be-shell. These phases are characterized by different spectral behaviours (a normal absorption-line spectrum or emission/shell spectrum) due to changes of electron density and/or of geometry of the CS envelopes (Doazan et al. 1988). There is no doubt that during the normal B phase, the electron density is too small to produce any considerable IR excess. According to Dougherty et al. (1994), for the IR excess in the L-band to be detectable, the density of the CS plasma has to be $\approx$ $10^{-11.6}~{\rm g\,cm}^{-3}$. To account for the maximum level of $p_{\ast}\approx0.5\ {\rm to}\ 1\%$ observed for the stars under peak 1 (with the above value of density and typical sizes of CS environment $\approx$ $10\ {\rm to}\ 15~R_{\ast}$), in the framework of the disk model, the half-opening angle of the disk cannot be smaller than 10$^\circ$ to 15$^\circ$. The important point is that in order to explain a bi-modal character of the distributions, one can suggest that the transition time between the normal B and Be/Be-shell phases is much shorter than the duration of the phases themselves. In this case, we can observe an object either at the normal B phase (with $E(V-L)\approx 0$) or at the Be/Be-shell phase with E(V-L) values from the distribution under peak 2. In principle, our suggestion is supported by the data of Doazan et al. (1986), who found significant spectral variations in $\theta$ CrB on a time scale of 3 months when the star changed its phase from shell to normal. At the same time, the duration of the phases in Be stars (or an envelope lifetime) is typically of the order of decades (Cramer et al. 1995; Fox 1991). Our suggestion is consistent also with the hypothesis that the formation of CS envelopes in Be stars is due to short lasting huge mass ejections (Hubert & Floquet 1998; Moujtahid et al. 1999; Zorec et al. 2000a, 2000b). All the above may be also valid to explain bi-modality in E(V-12) distributions (if it exists). Hubert-Delplace et al. (1981) and Mennickent et al. (1994) noted that the time scale of phase transitions is shorter in early-type Be stars than in late spectral types. Thus we would expect less pronounced bi-modality in late types too, this not being obvious in our data. The apparent absence of bi-modality for late spectral types (B4-B5 and B6-B9.5) can be explained by the notion that the second peak would be placed toward smaller values of E(V-L). Its possible location is perhaps masked by a combination of the chosen binwidth in association with the small statistical sample.

7 Distributions of intrinsic polarization for different subgroups of Be stars

As was noted before, (see Sect. 4.1 and Fig. 2b), Be stars of earlier spectral type tend to exhibit larger values of intrinsic polarization. Thanks to the large bank of calculated values of $p_{\ast }$ ($\approx$500 objects), it is possible to investigate the above mentioned behaviour in detail.

 

 

Table 6: Intrinsic and true polarization for different spectral subtypes and subgroups of Be stars

Spectral type:	O-B0	B1	B2	B3	B4	B5	B6	B7	B8	B9
No. of stars:	47	118	146	58	22	25	20	23	23	11
$\overline{p_{\ast}}$, (%)	0.83	0.85	0.75	0.59	0.52	0.48	0.47	0.45	0.29	0.22
$\overline{p}_{\rm t}$, (%)	1.25	1.28	1.13	0.88	0.78	0.72	0.70	0.67	0.43	0.33
Spectral range	O-B1.5		B2-B2.5	B3-B5.5			B6-B9.5				O-B9.5
No. of stars:	166		146	105			77				495
$\overline{p_{\ast}}$, (%); $\sigma_{\rm p}$	0.84   0.48		0.75   0.48	0.55   0.48			0.37   0.28				0.68   0.48
$\overline{p}_{\rm t}$, (%)	1.27		1.13	0.82			0.56				1.02

Analysis of the intrinsic polarization distribution for classical Be stars was first performed by McLean & Brown (1978) for a sample of 67 objects. They noted a rapid rise of the distribution toward smaller values of $\,p_{\ast}$. The distribution established here (corrected for errors) for a sample of 495 objects (see Fig. 8) is similar to that obtained by McLean & Brown (1978). However, with the chosen binwidth (0.3%), our distribution peaks at $p_{\ast}\approx 0.4\%$ with the average weighted value $\overline{p_{\ast}}= 0.68\%$ and $\sigma_{\rm p}=0.48\%$. Our large sample allows us also to study for the first time the distributions of $\,p_{\ast}\,$ for the representative samples of Be stars of different spectral subgroups. These distributions, also error corrected, are shown in Figs. 9a-c. For O-B1.5 stars we used a slightly larger binwidth (0.35%) because of the possibly larger uncertainties in determination of interstellar components. It is clearly seen that the histograms are significantly different for each spectral subgroup. In contrast to late type stars, the distributions for early spectral types (O-B1.5 and B2-B2.5) do not peak at the value $p_{\ast}\approx 0\%$. This result is significant in spite of the relatively large standard deviations (see Table 6). Similar to the analysis for E(V-L) distributions, we performed a more detailed study of the B3-B5 spectral subgroup and found again that such behaviour is present with just B3 stars, but it disappears for B4-B5 stars.

	Figure 8: Distribution of intrinsic polarization (corrected for errors) for the entire sample of Be stars. The vertical dashed line indicates the mean value of polarization
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	Figure 9: Distributions of intrinsic polarization (corrected for errors) for different spectral subgroups of Be stars. The vertical dashed line indicates the mean value of polarization
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A clear difference in the values of $\,\overline{p_{\ast}}\,$ for different spectral subtypes and subgroups is apparent in the sense that there is a decrease in mean values of intrinsic polarization toward late spectral types (see Table 6). Physically this fact has been discussed previously by McLean & Brown (1978) in terms of less dense and/or less oblate envelopes for later stars. Application of various tests indicates that differences in the mean values of intrinsic polarization for different subgroups of Be stars are statistically significant. McLean & Brown (1978) found that there are no objects among Be stars with the intrinsic polarization in excess 2%. Our data allow us to study this suggestion statistically. Only 5 stars out of 497 objects (i.e. about 1%) show $p_{\ast}>2$% and 5 stars additionally exhibit $p_{\ast}=2$% (all of them are early type stars B0-B2). Moreover, we conclude that about 95% of Be stars exhibit intrinsic polarization on the level $0\% < p_{\ast} < 1.5\%$. This result is very important because Waters & Marlborough (1992) noted that in a framework of the single-scattering approximation (and geometrically thin disks, $\Theta\approx15\hbox{$^\circ$ }$) it is difficult to obtain high levels of polarization above 2%. The observed level of polarization can be explained either by a disk of large optical depth and a small opening angle or by a wide disk (or an ellipsoidal envelope) of small optical depth. To explain the polarization up to 3%, multiple scattering is considered in the series of the papers of Wood et al. (1996a, 1996b, 1997) and early by Daniel (1980). Nevertheless, there is much evidence that CS envelopes are optically thin, at least up to $\lambda\approx~10~\mu$m (see, for example, Gehrz et al. 1974); this also follows from multicolour IR polarimetry of some Be stars by Jones (1979) (as the position angle does not change from the visual to 2.2 $\mu$m). From comparisons of position angles of CS disks derived from optical interferometry and intrinsic position angles of polarization for seven Be stars, Quirrenbach et al. (1997) reached the same conclusion. As follows from our analysis, the percentage of Be stars with $p_{\ast}>1.5$% is negligible and multiple scattering (and strongly flattened disks: $\Theta<5\hbox{$^\circ$ }$ and $\tau>1$) may occur only in a few objects. Note that the upper limit in intrinsic polarization can be larger if we consider the true polarization by removing the aspect angle effect, using the method suggested by McLean & Brown (1978): $\overline{p}_{\rm t}=\frac{3}{2}\overline{p_{\ast}}$. For reference, the parameters of true polarization for Be stars of different spectral subtypes and subgroups are also given in Table 6. Even in this case, however, the fraction of Be stars with $p_{\rm t}>2$% does not exceed 10%.

The next point to consider is the possible differences in polarization of Be and Be-shell stars. This comparison is quite important as there is active debate on the change in the structure of CS envelopes of Be and Be-shell stars (Marborough et al. 1993; Hanuschik 1996; van Kerkwijk et al. 1995). It was found earlier that for some objects, the level of polarization increases when a star goes through a shell phase (Arsenijevic et al. 1987). Using a small statistical sample, Fox (1993b) noted a little higher level of polarization of Be-shell stars (12 objects, $0.77\pm0.41\%$) in comparison with non-shell stars (21 objects, $0.42\pm0.39\%$) (note, however, the magnitudes of the errors). On the other hand, Ghosh et al. (1999) found no distinction in polarization level between 24 Be and 5 Be-shell stars. Using the data of Hanuschik (1996) we are able to identify 46 Be-shell stars which have shown shell characteristics at some point in their observed history and compare their behaviour with others. The average values of the parameters for non-shell stars have been calculated using the data from Tables 4 and 6 for the same specific population of stars of different spectral subtypes as in the shell subgroup. As follows from Table 7, no differences (in polarization and IR excesses) exist between Be non-shell and Be-shell stars. A compilation of polarimetric data for Pleione (see WUPPOL 1999) for different phases (Be and Be-shell) also shows the absence of a significant changes in the polarization level. This suggests that the phase transition Be $\Leftrightarrow$ Be-shell has no strong affect on the parameters describing the inner parts of the disks. Note, however, that all Be-shell stars show significant variability in photometry and polarimetry and the averaging over a long period is not a correct procedure. For this reason, the above discussion can be considered as only a preliminary suggestion. A statistical analysis of the rotational characteristics of Be-shell stars is much more informative as their $v\sin{i}$ values do not depend on the phase changes. In this context it is interesting to investigate an inclination aspect of Be and Be-shell stars as discussed earlier by Moujtahid et al. (1999). Although the errors are large, the mean viewing angle ("i'') of Be-shell stars derived from the ratio $\overline{v_{\rm t}}/\overline{v_{\rm c}}$ is larger on average than that for other Be stars ( $52\hbox{$^\circ$ }\pm20\hbox{$^\circ$ }$ and $40\hbox{$^\circ$ }\pm19$ respectively). Thus, Be-shell stars are viewed closer to the CS disk plane and none of them are observed at the angle less than $30\hbox{$^\circ$ }$.

 

 

Table 7: Differences in the observed characteristics of Be-shell and Be stars. The number of stars (N) used to calculate the average values of the parameters is indicated in parentheses

	$p_{\ast }$	E(V-L)	E(V-12)
Be-shell (N stars)	0.61$\%$(44)	$0\mbox{$\,.\!\!\!^{\rm m}$ }61$(40)	$1\mbox{$\,.\!\!\!^{\rm m}$ }49$(31)
non-shell (N stars)	0.57$\%$(448)	$0\mbox{$\,.\!\!\!^{\rm m}$ }47$(203)	$1\mbox{$\,.\!\!\!^{\rm m}$ }34$(119)

On the other hand, there is clear difference in average values of "i'' between Be-shell stars of different spectral subgroups:
O9-B1.5:   $42\hbox{$^\circ$ }\pm12\hbox{$^\circ$ }$       B2-B2.5:   $41\hbox{$^\circ$ }\pm10\hbox{$^\circ$ }$
B3-B5.5:   $59\hbox{$^\circ$ }\pm18\hbox{$^\circ$ }$       B6-B9.5:   $55\hbox{$^\circ$ }\pm17\hbox{$^\circ$ }$.
If the transition to the Be-shell phase is due to mass loss preferably close to a CS disk plane, the opening angle of CS disks may be smaller for late type stars.

The last point for consideration is a deviation of the distribution peaks from $p_{\ast}=0\%$ for early type stars. Although the problems in detection of the peak in a distribution with very small values of p are discussed by McLean & Brown (1978), the differences detected here are probably real. The first possibility to explain the behaviour is that most early type stars in our sample are more distant objects. The estimates of interstellar components for them are not so precise, so that the displacement from 0% can be attributed to inaccuracy of its removal. On the other hand, the behaviour can be explained if the suggestion on the randomly oriented rotation axes of envelopes is not true. This point will be discussed in a separate paper. Finally, it may be due to more complicated geometries associated with the CS environments of early-type stars.

8 Discussion

An excellent discussion on the problems of analysis of the polarization distribution of classical Be stars has been presented by McLean & Brown (1978), and a detailed discussion of near and far IR characteristics was made by Coté & Waters (1987) and Dougherty et al. (1991, 1994). Such kind of theoretical work is beyond the purpose of our paper and we will discuss some of our results only briefly. Our very large statistical sample allows us to explore possible correlations between the parameters discussed here and to make some suggestions on CS envelope geometries.

First, no clear correlations exist between the intrinsic polarization and near IR excesses neither for a sample of 209 stars (see Fig. 10a) nor for different spectral subgroups.

	Figure 10: Correlation between near a) and far b) IR excesses and intrinsic polarization for the entire sample of Be stars
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This negative result is very important because it indicates that the regions in CS envelopes which are responsible for generating the polarization and near IR excesses are different. Note that Coté & Waters (1987) and Waters & Marlborough (1992) found and discussed a possible correlation between the polarization and far IR excesses at 12 $\mu$m, but on the basis of very small sample (46 stars). To study this suggestion, we investigated the correlation between the values of polarization and the excesses at 12 $\mu$m (E(V-12)) for 125 Be stars. The data of E(V-12)were taken from Coté & Waters (1987) and Oudmaijer et al. (1992). A comparison of Figs. 10a and 10b indicates that, in contrast to near IR excesses, the relation between the intrinsic polarization and far IR excesses, with a factor three improvement of sample size, really exists in the form of a triangular shaped distribution similar to that found by Coté & Waters (1987). With a few exceptions (6 stars, three of them Be-shell), there are no stars with high polarization and low far IR excesses. The well defined upper limit is described by the empirical relation:

$${p}^{V}_{\ast}=0.65\times E(V-12).$$ (5)

Note that Coté & Waters (1987) plotted B band values of $p_{\ast }$ against E(V-12) and the upper limit for their dependence is described by ${p}^{B}_{\ast}=0.83\times E(V-12)$. The ratio of p^B/p^V, however, is virtually constant for all classical Be stars (see, for example, Ghosh et al. 1999) and is equal to $1.28\pm0.09$. Taking into account this ratio, we conclude that the upper limits for our dependence and for the dependence of Coté & Waters (1987) are exactly the same. Reasons for the lack of correlations between $p_{\ast }$ and E(V-L) in contrast to the $p_{\ast }$: E(V-12) triangular distribution is not clear. As noted by Poeckert & Marlborough (1976) and Coté & Waters (1987), the bulk of the polarization is produced in layers within <3 $R_{\ast}$, while the 12 $\mu$m excess can originate from layers that are further out. However, it was shown recently by Fox (1993b), that two parameters are important for continuum linear polarization in Be stars: the envelope oblateness and the location of the inner boundary of the cool extended envelope (which is typically $\approx$5 $R_{\ast}$). The behaviour seen in Fig. 10b does not contradict to the suggestion that $p_{\ast }$originates in those regions of CS envelopes which are optically thin. It may reflect a link between far IR excesses (or polarization) and the size of CS envelopes.

Note that although the correlation between near and far IR excesses clearly exists (with a correlation coefficient of $r=0.65\pm0.06$), there is a significant scattering of the data points (see Fig. 11).

	Figure 11: Correlation between near and far IR excesses for the entire sample of Be stars. The dashed lines represent a possible non-linear fit for the observational data for E(V-L): $E(V-m_{\lambda }$) dependencies where $m_{\lambda }$ is the magnitude at 12 $\mu$m, 25 $\mu$m and 60 $\mu$m respectively. The data points marked in this figure concern E(V-L) and E(V-12) data
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The large spread is difficult to explain if near and far IR excesses arise in the same gaseous disks and have the same nature (free-free and free-bound radiation). The correlation coefficients for different spectral subgroups show an increase toward late spectral type stars ( $r=0.36\pm0.14$, $0.55\pm0.14$ and $0.66\pm0.11$ for B2, B3-B5 and B6-B9 respectively). For O-B1.5 stars, $r=0.71\pm0.14$ but the statistical sample is too small (13 objects). Moreover, a visual inspection of Fig. 11 leads to the suggestion that the dependence E(V-L): E(V-12) is not linear and there is a more rapid increase of E(V-L) in the region of large E(V-12) values. The same conclusion follows from the consideration of the E(V-L): E(V-25) and E(V-L): E(V-60)dependencies. In contrast, the relations between the excesses at 12 $\mu$m, 25 $\mu$m and 60 $\mu$m are quite linear with correlation coefficients $r\approx0.8\ {\rm to}\ 0.9$.

On the one hand, the behaviour can be due to the fact that CS envelopes are optically thin at 3.5 $\mu$m but optically thick at $\lambda>12~\mu$m (Gehrz et al. 1974). This may lead to the behaviour indicated for E(V-L): $E(V-m_{\lambda }$) dependencies. The theoretical interpretation of these dependencies is beyond the purpose of our paper note, however, that a similar behaviour was discussed theoretically by Dougherty et al. (1994). The regions of large E(V-L) and E(V-12) values in Fig. 11 are occupied mainly by early type stars (see Fig. 2a and Coté & Waters 1987). These hotter stars can ionize larger volumes of CS matter, leading to the increase in E(V-L)excesses). Poeckert & Marborough (1982) noted that the dramatic change in ionization structure of CS envelopes is obvious around $T_{\rm e}\approx15\,000\ {\rm to}\ 20\,000$ K corresponding to spectral types B4-B6. At the same time the effective surface area of the emitting regions which are responsible for the excess E(V-12) for B0-B5 stars does not change strongly (see previous section). On the other hand, it is possible to suggest the presence of an additional component to CS envelopes for early spectral type objects close to the star which may be responsible for an increase of near IR excesses. The absence of this second component in late type stars leads to the increase of correlation coefficients for E(V-L): E(V-12) dependence from B2 to B9 stars.

	Figure 12: Correlation between near a) and far b) IR excesses with $v\sin{i}$values for the entire sample of Be stars
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	Figure 13: Correlation between near a) and far b) IR excesses with the ratio of $v\sin{i}/v_{\rm c}$ for the entire sample of Be stars
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Another important negative result is the absence of any correlation between near-far IR excesses and $\,v\sin{i}\,$ for 220 and 118 Be stars respectively (see Figs. 12a,b) (and for any spectral subgroup considered separately). This behaviour may result from the small optical depth of CS environment for the majority of Be stars in the near IR. Another possibility to account for this behaviour is a suggestion of either a more spherical geometry for the CS environment where the near IR emission arises or the presence of two independent regions - a disk component and a spherical component or polar lobes, both responsible for the near IR emission. The suggestion on two distinct regions in CS environment in classical Be stars (a diffuse polar stellar wind and a dense equatorial disk) is discussed currently in literature (see, for example, Porter, 1998 and references therein). Moreover, in context of polarization modeling, Fox (1991) supposed that for some Be stars the polar region can be the dominant source of free electrons. To minimize the uncertainties due to mixing in a same diagram involving stars with different masses and luminosity classes, we re-construct Fig. 12 using the ratio $v\sin{i}/v_{\rm c}$ instead of the values of $v\sin{i}$. However, as follows from Fig. 13, no correlations appear even with this adjustment. As was noted above, however, CS envelopes are optically thick at 12 $\mu$m. In the framework of this assumption, strongly flattened disks would produce high colour excesses for the case of pole-on orientation and significantly lower excesses at high inclinations (edge-on). For the geometrically thick disks (with $\Theta >50\hbox{$^\circ$ }$) we would expect an opposite behaviour, namely high excesses at $i=90\hbox{$^\circ$ }$ and lower excesses at $i=0\hbox{$^\circ$ }$. The absence of correlations between E(V-12) and $v\sin{i}$ indicates that the effective surface area of emitting regions is approximately equal for pole-on ( $S_{1}(i=0\hbox{$^\circ$ })\simeq \pi R^2$) and edge-on ( $S_{2}(i=90\hbox{$^\circ$ })\simeq 2\pi R^2\tan\Theta$) cases, where R is the radius of the disk and $\Theta$ is its half-opening angle. Regarding the disk-like structure, from simple geometrical considerations, this is realized for disks with a half-opening angle $\Theta\approx25\hbox{$^\circ$ }$. Such behaviour can never be obtained for ellipsoidal envelopes as the effective surface area (and thus E(V-12)) would be always larger for a pole-on orientation! In addition, note that according to Fox (1991, 1994) optically thin ellipsoidal envelopes are unable to produce polarizations in excess of 1.2%. Thus, the implication of this configuration of CS envelopes is at least questionable. Assuming $\sigma_{E(V-12)}\approx0\mbox{$\,.\!\!\!^{\rm m}$ }2$, for the difference in E(V-12) not to be detectable at 3$\sigma$ level, the ratio of the effective surface areas for pole-on and edge-on disks cannot exceed 0.6. With this ratio the range of the half-opening angle of CS disks is $10\hbox{$^\circ$ }<\Theta<40\hbox{$^\circ$ }$. Waters (1986) showed that for a disk with $\Theta=10\hbox{$^\circ$ }$, the effect of inclination on the excess flux cannot be negligible for $i>50\hbox{$^\circ$ }$. Our estimate of $\Theta$ is in agreement with that of Waters (1986) and Hanuschik (1996) but it clearly contradicts theoretical predictions for WCD models of Bjorkman & Cassinelli (1993) and Owocki et al. (1994). De Araújo et al. (1994) have determined the angle of a CS envelope where the density has decreased by a factor of 1/e with respect to the equatorial value as about $17\hbox{$.\!\!^\circ$ }5$ which is also in agreement with our data. Our large statistical sample allows us to suggest that any uncertainties due to non-simultaneous observations and/or variability of observed parameters should not affect the results significantly and the lack of any strong correlations is a real outcome.

However, if all Be stars are surrounded by CS disks, correlations might be expected between the intrinsic polarization and projected rotational velocities. This dependence was first investigated by McLean & Brown (1978) and recently for a larger sample of stars by Yudin (2000). In both cases, the relation between $p_{\ast }$ and $v\sin{i}$ shows a triangular distribution but with very different boundaries. According to McLean & Brown (1978), the polarization exhibits a maximum scatter for the largest values of $v\sin{i}$. In contrast, Yudin (2000) found that $\,p_{\ast}\,$ reaches maximum values for stars with projected rotational velocities in the intermediate range of 150 kms^-1 $< v\sin{i} < 250$ kms^-1 (Feature 1) and there are none (with a few exceptions) with large polarization among the stars with relatively low $v\sin{i}$ (<100 kms^-1), (Feature 2) and very high $v\sin{i}$ (>300 kms^-1), (Feature 3). Note that Yudin (2000) used the data of intrinsic polarization for about 90 objects and his result is well confirmed in the present study for the largest list of intrinsic polarization of classical Be stars currently available (337 objects, see Fig. 14). The distribution constructed here has a peak at the value $v\sin{i}\approx 200$ kms^-1 and the boundary lines are described by:

$p_{\ast}(\%) = 0.01(v\sin{i})$ for $0<v\sin{i}<200$ kms^-1and

$p_{\ast}(\%) = 4-0.01(v\sin{i})$ for $200<v\sin{i}<400$ kms^-1. The behaviour corresponding to Feature 2 has already been discussed in terms of high inclinations of the envelope symmetry axis to the line of sight (McLean & Brown 1978). While the polarization by scattering in any optically thin axisymmetric CS envelope is proportional to the optical depth and $\sin^{2}{i}$, the increased scatter in the $p_{\ast }$: $v\sin{i}$ diagram would be observed toward the larger values of $v\sin{i}$ (Feature 1). However, polarization of about 2% is the limit in the canonical model adopted for Be stars which is well confirmed in our study. Fox (1991), however, showed theoretically that the inclination dependence is more complicated than in the model of Brown & McLean (1977) and that the expression $p_{\ast}\propto\sin^2{i}$ no longer holds. Again, to minimize any uncertainties due to mixing of stars with different stellar masses and luminosity classes, we re-construct Fig. 14 using the ratio $v\sin{i}/v_{\rm c}$ instead of the values of $v\sin{i}$. The same triangular shaped distribution is apparent in Fig. 15. The peak of the distribution is at the value $v\sin{i}/v_{\rm c}=0.41$, corresponding to the value $v_{\rm t}/v_{\rm c}\approx0.52$. The boundary lines are described as follows:

$p_{\ast}(\%) = 4.88(v\sin{i}/v_{\rm c})$ for $0<(v\sin{i}/v_{\rm c})<0.4$ and

$p_{\ast}(\%) = 3.39-3.39(v\sin{i}/v_{\rm c})$ for $0.4<(v\sin{i}/v_{\rm c})<1$.

Finally, note that less than 3% of Be stars are located outside of the triangular distribution.

Let us consider now a few hypothesis in context to Feature 3. First, to account for this behaviour, it must be remembered that a rapidly rotating Be star (with the value of $v\sin{i}\approx 300$-400 kms^-1) is at the stability limit and has a large coefficient of non-sphericity.

	Figure 14: Correlation between $p_{\ast }$ and projected rotational velocities for the entire sample of classical Be stars
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	Figure 15: Correlation between $p_{\ast }$ and the ratio of $v\sin{i}/v_{\rm c}$ for the entire sample of classical Be stars
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It is well known that the rapid rotation of the star produces two effects, namely flattening of the star at the poles and, secondly, a gravitationally induced darkening of the equatorial region. Both of these effects contribute to a polarization whose direction is perpendicular to the polar axis i.e., orthogonal to the vector of polarization arising in an optically thin CS gaseous disk (Bjorkman & Bjorkman 1994). Consequently, the observed polarization, being the vector sum of the two components (CS polarization and star's photospheric polarization), tends to diminish toward the highest $v\sin{i}$ values, beginning at some limiting value, this leading to the appearance of the "triangular" distribution of $p_{\ast }$: $v\sin{i}$. Note however, that the depolarizing effect is small, ( $\Delta p_{\ast}<0.3\%$, see Bjorkman & Bjorkman 1994), especially taking into account that the decrease of $\,p_{\ast}\,$ is clearly obvious from the relatively small value $v_{\rm t}/v_{\rm c}\approx0.6$ (see further discussion). Moreover, scattering of stellar polarized radiation in a CS envelope would lead to appearance of circular polarization (also true for multiple scattering) but this is not detected in Be stars at the moment.

The second possibility to account for the above mentioned behaviour is that the triangular shaped distribution might be the result of a combination of two facts established above namely: larger mean projected rotational velocities and smaller absolute values of polarization for late spectral type stars. In this case, most late spectral type stars will be located at the lower right side of the diagram that tending to the decrease of polarization toward larger values of $v\sin{i}$ for the sample of Be stars. However, a similar diagram constructed with just B6-B9.5 stars again shows the same behaviour with a triangular distribution - as well as for the other spectral subgroups of Be stars treated separately. Therefore, the common behaviour is attributed to rapid rotation independently on spectral subtypes.

An alternative suggestion is that the stars with a rapid rotation may form narrower CS gaseous disks. Because the generated polarization is proportional to the ratio of thickness to length of the disk (or to an opening angle, see, for example, Dolginov et al. 1995), such narrower disks produce a significantly lower level of polarization due to increased dilution by unpolarized starlight. Inasmuch as we exclude from consideration multiple scattering, ellipsoidal envelopes, geometrically thick and strongly flattened disks, let us discuss now the optically thin CS disks with $10\hbox{$^\circ$ }<\Theta<40\hbox{$^\circ$ }$. For disk-like envelopes the maximum degree of polarization ( $p_{\ast}\approx1.5$-$2\%$ for the optical depth $\tau\approx0.5$) occurs at $\Theta\approx35\hbox{$^\circ$ }$ (Waters & Malborough 1992; Fox 1993a, - for single-scattering, plus attenuation) that is similar to the above derived upper limit of $\Theta$. As follows from the theoretical calculations of Fox (1993b, 1994) and Waters & Malborough (1992), a decrease of $\Theta$ from $35\hbox{$^\circ$ }$ to $10\hbox{$^\circ$ }$ leads to the decreasing in polarization from 1.5-2% to 0.5-0.7%. In addition, the above mentioned authors showed that polarization reaches a maximum at the inclination angle $i\sim85\hbox{$^\circ$ }$ rather than at $i\sim90\hbox{$^\circ$ }$. Thus, for exactly edge-on orientation of CS envelopes the observed polarization is reduced by about 0.3-0.5%. Taking into account the star's non-sphericity, a combination of these three effects can reduce the intrinsic polarization to zero for stars with rapid rotation, viewed edge-on. Although the models of disk formation around early-type stars currently discussed in the literature (rotation induced bi-stability and wind compressed disks) cannot explain the existence of CS disks with $\Theta>5\hbox{$^\circ$ }$, in both models the opening angle of a disk decreases with the increase of the ratio $\overline{v_{\rm t}}/v_{\rm c}$.

An interesting point is that the peak in the distribution $p_{\ast }$: $v\sin{i}/v_{\rm c}$ is different for our spectral subgroups (see Table 8). In the framework of the above suggestion, the compression of CS disks which affects the depolarization, sets in for lower values of $\overline{v_{\rm t}}/v_{\rm c}$ for early type stars. However, this hypothesis requires a more detailed theoretical study.

 

 

Table 8: Location of the peaks in $p_{\ast }$: $v\sin{i}/v_{\rm c}$, $p_{\ast }$: $v_{\rm t}/v_{\rm c}$ distributions and the average values of $\overline{v\sin{i}}/\overline{v_{\rm c}}$, $\overline{v_{\rm t}}/\overline{v_{\rm c}}$for different spectral subgroups of Be stars

	O-B1	B2	B3-B5	B6-B9	O-B9
$(v\sin{i}/\overline{v_{\rm c}})_{\rm peak}$	0.33	0.41	0.47	0.56	0.41
$\overline{v\sin{i}}/v_{\rm c}$	0.40	0.45	0.56	0.62	0.50
$(v_{\rm t}/\overline{v_{\rm c}})_{\rm peak}$	0.42	0.52	0.60	0.71	0.52
$\overline{v_{\rm t}}/v_{\rm c}$	0.51	0.57	0.72	0.80	0.64

Finally, note that our study is essentially just one of statistical analysis only of some of the observational characteristics of Be stars and most physical explanations are only a simple sketch. To construct a self-consistent model of Be phenomena, a detailed theoretical interpretation of the new results presented here and earlier results well confirmed in the present study now needs to be undertaken.

9 Conclusions

A statistical analysis of some observational characteristics (polarization, projected rotational velocities and near-far IR excesses) for the largest currently available list of classical Be stars (627 objects) has been completed. Our main conclusions may be summarized as follows:

1.

Polarization and near IR excesses of Be stars plotted against spectral types show the same behaviour, namely a maximum mean value for B1-B2 spectral types with a decrease toward late spectral types;

2.

Projected rotational velocities for the different spectral subgroups of Be stars are statistically significantly different. We confirmed that for luminosity class V, late spectral type stars rotate faster on average and the ratio of $\overline{v\sin{i}}/v_{\rm c}$ for them is larger than for early spectral types. For the III-IV luminosity classes, no significant differences exist between $\overline{v\sin{i}}$ but the ratio of $\overline{v\sin{i}}/v_{\rm c}$ is also larger for late spectral type subgroups. All the above confirm well the results of previous studies;

3.

Analysis of the data allows us to conclude that for E(V-L) excesses, there is a bi-modal distribution for O-B1.5, B2-B2.5 and B3-B5.5 spectral subgroups. Such a behaviour does not appear for late spectral types (B6-B9.5). The second peak in E(V-L) distributions transits to the lower values from early to late spectral type subgroups;

4.

We conclude that about 95% of Be stars exhibit intrinsic polarization on the level $0\% < p_{\ast} < 1.5\%$;

5.

No correlation is found between $p_{\ast }$ and E(V-L). In contrast, a clear relation between $p_{\ast }$ and far IR excesses at 12 $\mu$m is confirmed;

6.

No correlation is observed between E(V-L), E(V-12) and $\overline{v\sin{i}}$ or $\overline{v\sin{i}}/v_{\rm c}$;

7.

A clear relation between the intrinsic polarization and $v\sin{i}$(as well as for $p_{\ast }$- $\overline{v\sin{i}}/v_{\rm c}$) is found in the form of "triangular distribution'' with a decrease of $p_{\ast }$ toward very small and very large $v\sin{i}$ (and $\overline{v\sin{i}}/v_{\rm c}$) values. The peak of this distribution is at values $\overline{v\sin{i}}\approx200{-}220$ kms^-1 and $\overline{v\sin{i}}/v_{\rm c}\approx0.41$( $\equiv\overline{v_{\rm t}}/v_{\rm c}\approx0.52$). The decrease of intrinsic polarization toward larger values of projected rotational velocities can be understood in the context of the larger oblateness of CS disks for the stars with a rapid rotation;

8.

From the analysis of correlations between different observational parameters we conclude that circumstellar envelopes for the majority of Be stars are optically thin disks with the range of the half-opening angle of $10\hbox{$^\circ$ }<\Theta<40\hbox{$^\circ$ }$ and the opening angle of CS disks may be smaller for late type stars.

Acknowledgements

I wish to express my sincere thanks to Dr. D. Clarke (Glasgow University) for useful discussions and for the help in the improving of the text. I also thank the referee (Dr. J. Zorec) for his helpful comments. We acknowledge the use of the SIMBAD database and this research has made use of the VizieR Service at Centre de Données Astronomiques de Strasbourg. The research described in this paper was made possible in part by the 99-02-16336 grant of RFBR.

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Online Material

Appendix I. Observational data of classical Be stars

HD	HR	MWC	Sp	E(V-L)	$p_{\rm intr}$	$\overline{v\sin{i}}$	E(V-12)	other name	remarks
144	7	2	B9IIIe	0.05	0.30	$146\pm14$		10 Cas
4180	193	8	B2Ve	0.87	0.23	$221\pm~9$	1.48	22 Cas,o Cas
5394	264	9	B0IVe+sh	1.76	0.43	$253\pm24$	2.13	$\gamma$ Cas
6811	335	420	B6.5IIIe	0.15	0.14	$~88\pm13$	0.44	$\phi$ And
7636		12	B2IIIne	$\approx$0.5	0.51	$148\pm~3$	1.86	V764 Cas
9709		426	B8V			$297\pm30$
10144	472		B3Vep	0.20	0	$207\pm~9$	-0.05	$\alpha$ Eri
10516	496	16	B0.5IVe+sh	1.80	2.17	$395\pm29$	2.36	$\phi$ Per
11415	542		B3V	0.12	0.28	$~38\pm~9$		$\varepsilon$ Cas
12302		23	B1Vep	0.55	0.80	$208\pm24$		V780 Cas
12856		25	B0pe		0.91	$163\pm~9$
13051		27	B1IVe	1.72	0.92	$157\pm~7$		V351 Per
13429		439	B3V			$255\pm26$
13661		29	B2Ve	0.5		$196\pm21$		V594 Per
13669		707	B2Vne			$304\pm22$
13758			B1V		1.48	$202\pm20$
13867		442	B5Ve			$~93\pm16$
13890		443	B1IIIe		0.95	$179\pm~8$
14422		37	B1Vpe	0.16	0.47	$~98\pm10$		V424 Per
14605		44	B1IVpe	0.70	0.47	$150\pm~3$		V361 Per
14850			B7III/IVe			$200\pm20$
15238		47	B3V	$\approx$0.15		$273\pm27$		V529 Cas
15450		48	B1IIIe		0.97	$242\pm24$		V555 Per
18552	894	455	B8Vne	$\approx$0.2	0.38	$260\pm12$
18877		721	B9Ve		0.45	$178\pm~1$
19243			B1Ve		0.78	$138\pm14$	1.91
20017			B7Ve		0.62	$249\pm46$
20134		64	B2.5Vep		0.90	$~87\pm~9$
20336	985	65	B2.5Ven	0.80	0.65	$286\pm19$	0.86	BK Cam
20340		723	B3Ve			$250\pm25$
20418	989		B5V		0.13	$262\pm~4$		31 Per
20899		725	B9e			$162\pm14$
21212		67	B1e	1.82	2	$106\pm18$	1.89
21362	1037		B6Vn		0.1	$313\pm21$
21455	1047		B7V		0.10	$138\pm~3$	1.02
21551	1051		B7V	0	0.85	$296\pm~8$
21641		727	B8.5Ve		0.26	$170\pm~1$
21650		68	B6Ve		0.6	$226\pm22$
22192	1087	69	B4.5Ve+sh	0.58	1.7	$297\pm41$	1.84	$\psi$ Per,37 Per
22298		70	B2Vne		1.12	$154\pm15$	1.25	CT Cam
22780	1113	463	B7Vne	0.38	0.47	$283\pm~8$
23016	1126		B8.5Vne	0.04	0.55	$278\pm14$		13 Tau
23302	1142	72	B6IIIe	0.13	0.09	$189\pm~7$	1.25	17 Tau
23478			B3IVe:		1.3	$170\pm~7$
23480	1156	73	B6.5IVe	0.55	0.18	$252\pm28$	0.57	23 Tau,V971 Tau
23552	1160	464	B8Vne	$\approx$-0.15	0.16	$218\pm~3$	0.73
23630	1165	74	B7IIIe	0.19	0.12	$173\pm25$	0.60	$\eta$ Tau,25 Tau
23800			B1.5IVe		1.26	$178\pm~5$
23862	1180	75	B8Vpe+sh	0.58	0.38	$295\pm15$	1.02	28 Tau
23982		76	B4e			$173\pm~5$	1.83
24479	1204	77	B9.5Ve	$\approx$0.2	0.21	$110\pm~9$	-0.04
24534	1209		O9.5Vep	0.50	0.48	$162\pm23$	0.89	X Per
25348		80	B1Vpnne		1.16	$202\pm20$		DE Cam
25799			B3V		0.45	$322\pm32$		V490 Per
25940	1273	81	B4Ve	0.57	0.17	$198\pm31$	1.53	48 Per,MX Per
26356	1289		B5Ve		0.06	$274\pm24$	1.07
26398		468	B7III			$167\pm~2$	1.57
26420		82	B3Ve		0.66	$186\pm29$
26670	1305		B6V	$\approx$1.1		$263\pm26$
26906		83	B2Ve		1.22	$230\pm11$		V586 Per
27846			B1.5V			$134\pm15$
28497	1423	86	B1Ve	1.21	1.02	$256\pm38$	2.26	DU Eri	Vega-type?
29441		733	B2.5Vne			$311\pm40$		V1150 Tau
29866	1500	88	B7.5IVne	0.28	1.19	$250\pm11$	0.80
30076	1508	89	B2Ve	0.84	0.17	$213\pm25$	1.04	DX Eri,56 Eri
30123		736	B8IIIe			$~90\pm~1$
32190			B1Ve		0.08	$130\pm13$		V1153 Tau
32343	1622	96	B3Ve	0.99	0.12	$118\pm11$	2.06	11 Cam,BV Cam
32990	1659		B2Ve	0.50	0.69	$~94\pm~2$		103 Tau
32991	1660	98	B3Ve	0.80	0.40	$198\pm~2$	1.88	105 Tau,V1155 Tau
33152		99	B1Ve		0.75	$~98\pm10$		V412 Aur
33232		100	B2Vne	0.75	0.92	$214\pm44$
33328	1679		B2IVne	0.26	0.24	$271\pm26$		$\lambda$ Eri
33461		101	B2Vnne	1.85	0.64	$145\pm61$	1.93	V415 Aur
33599			B5pe		0.7
33604		103	B2Ve		0.44	$134\pm20$
33988		104	B0Ve		0.22	$134\pm41$		12 Aur
34257		489	B5/8e			$218\pm20$
34921		107	B0IVpe	1.52	2	$234\pm~4$	1.94
34959	1761		B5Vpn		0.7	$233\pm45$		V1369 Ori
35165	1772		B5IVe+sh	0.12	0.7	$272\pm27$
35345		109	B1.5Vpe		0.65	$110\pm11$
35347		494	B1Ve		1.11$^\dagger$
35411	1788		B1V+B2V	0.15	0.20	$~58\pm~2$	0.15	$\eta$ Ori
35439	1789	110	B1Vpe	0.47	0.46	$266\pm23$	2.06	25 Ori, V1086 Ori
36012		757	B4Ve		0.10	$162\pm~3$		V1372 Ori
36376		760	B3/8e			$~58\pm~4$		V1374 Ori
36408	1847		B7IIIe			$~66\pm~6$	0.54
36576	1858	111	B1.5IVe	0.58	1.17	$231\pm13$	2.00	120 Tau
36665		501	B0/8e			$194\pm~1$
37115		114	B6Ve	0.86	0.37	$226\pm23$			Vega-type?
37202	1910	115	B1IVe+sh	0.96	1.13	$241\pm27$	1.83	123 Tau,$\zeta$ Tau
37330		116	B6V		0.26	$134\pm~7$
37490	1934	117	B3IIIe	1.02	0.27	$171\pm18$	1.07	$\omega$ Ori,47 Ori
37541		769	B5e			$154\pm12$
37657		118	B3Ve		1.32	$216\pm20$		V434 Aur
37795	1956	119	B7.5Ve	0.28	0.15	$184\pm~5$	0.86	$\alpha$ Col
37967	1961		B2.5Ve	$\approx$0.5	1.1		1.76	V731 Tau
38010		124	B1Vpe		0.33	$338\pm~3$	2.21	V1165 Tau
38087			B5Ve	0.49	2.5		3.79
38191		125	B1Vne		0.5
38708			B3Ve	$\approx$0.25	1.06	$226\pm~4$		V438 Aur
39018		779	B9e			$194\pm14$
39340			B3Ve		0.13	$222\pm22$
39478		129	B2Ve		2

Appendix I. Observational data of classical Be stars

HD	HR	MWC	Sp	E(V-L)	$p_{\rm intr}$	$\overline{v\sin{i}}$	E(V-12)	other name	remarks
40193			B8Vne		0.4
40978		131	B3Ve		0.58	$202\pm~3$		V447 Aur
41335	2142	133	B2Vne	1.57	0.60	$352\pm60$	2.23	V696 Mon
42054	2170	134	B4IVnne		0	$220\pm22$	1.72		Vega-type?
42529		794	B9e			$242\pm13$
42545	2198		B5Vne	$\approx$0	0.1	$273\pm23$		69 Ori
43285	2231	136	B5.5Ve	0	0.08	$260\pm12$
43544	2249		B2.5Ven	0.13	0.05	$303\pm~4$
43703		799	B1IVe		1.3	$216\pm~2$
44458	2284	138	B1Ve+sh	1.40	0.72	$224\pm21$	2.24	FR CMa
44506	2288		B3Ve+sh		0.39	$218\pm40$
44637		139	B2Vpe		0.6	$130\pm13$
44996	2309	526	B2.5Ve	0.12		$~98\pm33$
45314		140	B0IVpe	0.63	0.2	$245\pm25$
45542	2343	141	B6IIIe+sh	0.07	0.10	$193\pm14$	0.54	18 Gem,$\nu$ Gem
45626			B7pe	0.5	1.25
45725	2356	143	B3Ve+sh	1.04	0.24	$308\pm27$	2.25	$\beta$ Mon A,11 Mon
45726	2357		B3Vpe		0.13	$295\pm30$		$\beta$ Mon B
45871	2364		B5Ve	$\approx$0.4		$307\pm44$
45910		145	B2IIIpe+sh	1.85	0.70	$205\pm155$	2.23	AX Mon
45995	2370	146	B2Vne	1.20	0.91	$254\pm~5$	2.28
46380		530	B1Ve	$\approx$0.45	1.28	$262\pm12$		V728 Mon
46680			B9IV			$200\pm20$		$\mu$ Pic
47054	2418	150	B8Ve	$\approx$0.6	0.34	$226\pm20$	0.83
47359			B0.5Vp		0.7
47761		531	B1Ve		0.6	$~82\pm~5$		V733 Mon
48282			B3III		1.0
48917	2492	152	B2Ve	0.95	0.51	$200\pm20$	1.80	10 CMa,FT CMa	Vega-type?
49131	2501		B2IIIe			$295\pm30$	2.09	HP CMa	Vega-type?
49330		821	B0pnne		1.14			V739 Mon
49336	2510		B4Vne		1.65			V339 Pup	Vega-type?
49787		153	B2Ve		0.7	$186\pm~2$
49888		534	B3Ve		0.50	$250\pm~2$
49977		154	B2Ve	$\approx$0.5	1.95	$218\pm~4$		KU CMa
50013	2538	155	B1.5IVne	1.84	0.3	$201\pm24$	2.71	$\kappa$ CMa
50083		156	B2Ve		1.25	$178\pm~2$	2.27	V742 Mon
50123	2545	157	B6IVe	1.42		$226\pm~5$	1.84		Vega-type?
50658	2568	537	B8IIIe	0.19	0.6	$234\pm~5$		19 Aur
50696			B1Vnne		0.6
50820	2577	827	B3IVe	1.52		$119\pm12$	1.41
50938		540	B3Ve			$162\pm~5$		LL CMa
51354		160	B3Ve		1.0	$306\pm25$
51480		161	B8e	1.53	0.55		1.97	V644 Mon	young?
52244		163	B2IIIpne		1.33			LP CMa
52356			B4Ve			$394\pm58$		LQ CMa
52437	2628		B2Vne		0.12	$178\pm18$		FU CMa
52721			B2Ve	0.71	1.4	$243\pm93$	2.79	GU CMa
52812			B3V		1.05
52918	2648		B1Ve	0.35	0.4	$280\pm~7$	1.42	19 Mon
53257	2659		B9IV	$\approx$0.4		$304\pm~5$		44 Gem
53367		166	B0IVe	1.25	0.75	$~86\pm20$	2.53	V750 Mon
53416			B8e			$194\pm~2$
54309	2690	167	B1.5IVe	1.14	0.55	$210\pm30$	1.76	FV CMa
54464		836	B2.5Ve		0	$154\pm~1$
54475	2691		B2V			$200\pm20$
54786			B0pe		1
55135		168	B4Vne		0.30	$258\pm25$
55538			B2Ve		0.65	$146\pm11$		HI CMa
55606			B1Vpnne		0.55
56014	2745	170	B3IIIp+sh	0.33	0.25	$144\pm~5$	1.11	27 CMa,EW CMa
56139	2749	171	B2.5Ve	0.92	0.14	$109\pm~8$	1.15	$\omega$ CMa,28 CMa
56806			B1.5Vpnne		0.3
57150	2787	173	B2Ve	1.04	0.77	$254\pm35$	1.98	NV Pup	Vega-type?
57219	2790		B2Vne	0.10	0.03	$117\pm16$	0.26	NW Pup
57386		174	B2Ve		0.18	$274\pm15$
57393		556	B2Vnne		0.6
57539			B5III		0.5
58011			B1Vnep		0.1	$~18\pm~2$
58050	2817	176	B2Ve	0.69	1.12	$123\pm~6$		OT Gem
58127			B6e			$234\pm10$
58155	2819		B3V			$250\pm25$
58343	2825	177	B3Ve	0.86	0.40	$~57\pm15$	1.64	FW CMa
58715	2845	178	B8Ve	0.24	0.14	$244\pm~6$	0.60	$\beta$ CMi,3 CMi
58978	2855	179	B0.5IVne+sh	1.10	0.85	$245\pm32$	1.88	FY CMa
59094		561	B1Ve	1.18	0.18	$218\pm18$		V370 Pup
59498			B2Vne		0.12
60307			B3.5IIIe			$220\pm22$
60606	2911	183	B3Vne	1.21		$270\pm39$	2.17	OW Pup,Z Pup
60757		848	B5e			$170\pm~3$		V763 Mon
60848		184	O8Vpe	1.20	0.22	$247\pm31$		BN Gem
60855	2921	565	B2Ve	0.14	0.31	$239\pm~6$	1.03
61224	2932		B9IIIe			$234\pm~2$
61355			B3Vep		0.8
61925	2968		B5IV			$200\pm20$
62367		567	B8e			$114\pm~6$
62532			B1Vpnne		0.55
62654			B3Vne		0.7
62729		572	B3Ve			$194\pm~5$
62753		185	B2.5Vne		0.65			V387 Pup
62780		573	B2IIIne		1.90
63150			B0.5IVnep		1.5
63359			B1Vpe		0.9
63462	3034	186	B1IVe	1.21	0.2	$323\pm12$	1.71	o Pup
64109		187	B8e			$250\pm41$
64298			B2IVnne		0.68
65663			B8V	$\approx$0.30
65719			B2.5IIIe		0.15			V408 Pup
65875	3135	190	B2.5Ve	1.08	0.15	$153\pm19$	1.82	V695 Mon
66194	3147		B2.5Ve	0.24	0.32	$224\pm24$	0.67	V374 Car	Vega-type?
68468		577	B4e	$\approx$0.5		$122\pm~8$
68980	3237	192	B1.5Ve	0.79		$127\pm12$	1.97	r Pup,MX Pup
69168			B3V		0.16
69404		193	B2Vne		0.4	$298\pm30$
69425			B1Vne		1.11
71072		857	B3.5IVe	$\approx$0		$202\pm~2$

Appendix I. Observational data of classical Be stars

HD	HR	MWC	Sp	E(V-L)	$p_{\rm intr}$	$\overline{v\sin{i}}$	E(V-12)	other name	remarks
71510	3330		B2Ve	0.06		$162\pm18$
71934			B2IVne		0.46
72014			B1Vnne		0.65	$221\pm~8$
72063		858	B2Vne		0.9
72067	3356		B2Ven		0.2	$150\pm15$	1.54
74401			B1.5IIIne		1.3
75081	3488		B9Ve			$223\pm22$
75311	3498	195	B3Vne	0.18		$269\pm24$	0.65	V344 Car
75658			B1IVe		0.45
76534			B2V	0.29	0.53	$202\pm20$		OU Vel
77032			B2Vne		1.14
77320	3593		B2.5Ve	0.83	1.24	$284\pm38$		IU Vel
78764	3642	196	B2IVe		0.08	$127\pm~6$		V345 Car
79351	3659		B2.5V			$~35\pm~5$		a Car
79621	3670		B9V			$170\pm17$
79778			B2IVne		0.2
80834			B5Vnne		0.73
81357		859	B8.5e	$\approx$1.5		$162\pm16$			binary
81654			B2Ve		0.69
81753	3745		B5Ve			$300\pm30$			Vega-type?
82830			B1IIIne		2.0
83043			B2Vpne		0.92	$174\pm17$
83060			B2Vnep		0.50
83597			B1.5Ve		1.78
83953	3858	197	B6Ve	0.46	0.32	$273\pm18$	1.78	I Hya	Vega-type?
84523			B3Ve		0.25
84567	3878		B0.5IIIne		0.25	$242\pm24$
86272			B2Vne		0.5
86612	3946	581	B4Ve	0.82	0.01	$187\pm19$		OY Hya
86689			B2Vne		0.71
87380			B1IVn		0.20	$230\pm23$
87543	3971		B7IVne		0.20		1.44
87901	3982		B7V	0.12	0.1	$258\pm32$	0.11	$\alpha$ Leo
88661	4009	199	B2IVpne	0.97		$262\pm40$	1.89	QY Car
89080	4037		B8IIIe+sh	0.34	0.45	$207\pm13$	0.71	$\omega$ Car
89884		582	B5IIIe	$\approx$0.3		$256\pm33$
89890	4074	201	B3IIIe			$~70\pm~7$	0.27
90187			B1IIIne		0.45
90563			B2Ve		0.9
91120	4123	205	B9Vne+sh	0.13	0.23	$265\pm33$		Z Crv,5 Crv
91188			B4IIIe		1.25
91465	4140	208	B4Ve+sh	0.98	1.16	$255\pm~5$	1.64	P Car
92027			B2.5Vne		0.71
92061			B1.5IIIe		0.05	$124\pm12$
92714			B2Ve		1.1
92938	4196		B4V		0.24	$130\pm~3$		V518 Car
93030	4199		B0Vp	0.3	0.02	$145\pm~6$		$\theta$ Car
93237	4206		B5IV		0.25	$~70\pm~7$		DR Cha
93561			B1IIInne		1.0
93563	4221		B8.5IIIe+sh			$237\pm37$	0.91
93683			B1IVnep		1.8
94648			B2Ve	$\approx$1.55	0.3			FK Car
96357			B1.5Vne		1.3
96864			B1.5IVnep		1.4
98624			B1Vne	$\approx$1.3	1.1
98927			B1IIInne		0.8	$102\pm10$
99354			B1IIIne		0.84
100199			B1III		0.5	$187\pm19$
100324			B2Vne		2	$199\pm20$		KP Mus
100673	4460		B9Ve			$132\pm~5$	0.12
100856			B2IVpe			$175\pm18$		V843 Cen
100889	4468		B9.5V		0.04	$182\pm~9$	0.05	$\theta$ Crt,21 Crt
102742			B3Ve		0.15
102776	4537		B3Vne			$222\pm12$	0.13
105382	4618		B6IIIe	0		$111\pm10$	-0.02	V863 Cen
105435	4621	219	B2IVne	0.97	0.33	$207\pm46$	1.91	$\delta$ Cen
105521	4625		B3IVe		0.38	$159\pm40$		V817 Cen
105675			B2IVne		1.55
105753			B2IVne		0.7
107348	4696	221	B8Vne	0.27	0.04	$229\pm24$		$\zeta$ Crv
109387	4787	222	B6IIIpe+sh	0.63	0.24	$200\pm15$	1.73	$\kappa$ Dra,5 Dra
109857			B8Ve	$\approx$0.25			0.96		Vega-type?
110335	4823	223	B6IVe+sh		1.05	$208\pm34$	1.44	39 Cru,CH Cru
110432	4830	224	B2Vpe	1.29	1.3	$277\pm56$	2.13	BZ Cru
110863			B2IIInep		1.94	$182\pm18$		DQ Cru
112078	4897		B4Vne		0.07	$283\pm11$	-0.11	$\lambda$ Cru
112091	4899	225	B5Vne	0.62	0.12	$209\pm18$	1.57	$\mu^{2}$ Cru
112999			B6IIIn		0.9
113120	4930	226	B1.5IIIne	0.93	0.4	$307\pm~4$	1.82
113605			B2Vpe		0.7
113902	4951		B8V		0.04	$261\pm26$
114200			B1IIIne		0.4
114800			B2IVne		1.7	$204\pm~3$		V955 Cen
114981			B4Vne		0.02			V958 Cen	Vega-type?
116781			B0IIIne		0.8
117111			B2Vne		0.6	$186\pm19$
117357			B0.5IIIne		0.8	$~78\pm~8$
118246			B3Ve		0.05	$277\pm18$		GP Vir
120324	5193	229	B2IV-Ve	0.97	0	$159\pm~9$	0.10	$\mu$ Cen	Vega-type?
120991	5223	230	B2IIIep	1.22	0.09	$~78\pm~7$		V767 Cen
121556			B2IVnep		1.73
121847	5250		B8Ve+sh	-0.12	0.04	$193\pm~7$	-0.11	47 Hya
122450			B2IIIe		0.66
122669			B0.5Ve	1.1	0.48
122691			B1Vnne	1.3	1.2
124367	5316	231	B4Vne	0.82	0.35	$264\pm36$	2.07	V795 Cen
124771	5336		B3V		0.37	$201\pm~8$	0.61	$\epsilon$ Aps	Vega-type?
125159			B2Ve		0.3
127489			B2Ve		0.46	$186\pm19$		V1010 Cen
127617			B7III/IVe			$152\pm13$
127972	5440	232	B1.5Vne+sh	0.19	0.47	$291\pm32$	0.83	$\eta$ Cen	Vega-type?
129954	5500		B2.5Ve		0.67	$180\pm18$		CO Cir
131492	5551		B3Ve	0.68	0.5	$153\pm38$	0.49	$\theta$ Cir
133738			B1.5Ve		0.81	$161\pm~3$
134401			B2Vne		1.2

Appendix I. Observational data of classical Be stars

HD	HR	MWC	Sp	E(V-L)	$p_{\rm intr}$	$\overline{v\sin{i}}$	E(V-12)	other name	remarks
134481	5646		B9Ve	0.05		$172\pm~8$	0.29	$\kappa^{1}$ Lup
135160	5661	234	B0.5Ve	0.27	0.4	$155\pm19$
135734	5683		B8Ve			$243\pm92$	0.62	$\mu$ Lup
136415	5704		B5IVe+F8		0.36	$273\pm30$		$\gamma$ Cir
136556			B1Vne		1.45
137387	5730	235	B3IVpe+sh	1.02	0.45	$299\pm51$	1.43	$\kappa^{1}$ Aps
137432	5736		B4Ve	$\approx0$	0.08	$127\pm~7$
137518			B1IIInep		1.0
138749	5778	237	B6Vnne	0.17	0.15	$327\pm22$	-0.05	4 CrB,$\theta$ CrB
138769	5781		B3IV		0.05	$~93\pm16$		KT Lup
139431			B2Ve			$267\pm32$
139790			B3IIIne		1.8
140336			B1IIIne		0.6
140784	5860		B8Vn		0.1	$357\pm36$
140926			B2IVne		0.42
140946			B2IVne		0.5
141637	5885		B1.5Vn		0.62	$266\pm34$		1 Sco
141926			B2IIInne		1.05		2.81
142184	5907		B2Vne	0.20	0.17	$315\pm41$	0.76
142237			B3IVne		0.65
142926	5938	584	B9pe+sh	0.29	0.16	$298\pm27$		4 Her,V839 Her
142983	5941	239	B3IVe+sh	0.62	0.73	$369\pm22$	1.64	FX Lib,48 Lib
144320			B1Vne		0.85	$186\pm19$		V364 Nor
144708	6002		B9V			$265\pm32$
144965			B3Ve	$\approx$0.5
144970			B0Ve		1.0	$138\pm14$
145846			B1.5Vne		0.85	$142\pm27$
146444			B4Ve		1.25	$268\pm34$
146463			B2IVnep		1.5
147274			B7IVe		0.65
147302			B3Vne		0.35	$369\pm53$
147756			B1.5Vne		0.1
148184	6118	241	B1.5Vpe	1.10	0.41	$139\pm14$	2.18	7 Oph,$\chi$ Oph	Vega-type?
148259			B2Ve		0.3	$~86\pm~9$		OZ Nor
148877			B7Vep		0.25
149298			B1Ve		1.3	$150\pm15$
149313			B2IVne		0.39
149671	6172		B7IVe		0.16	$230\pm23$		$\eta^{1}$ TrA
149729			B2Vn		0.38	$258\pm26$
149757	6175		O9Ve	0.50	0.82	$341\pm32$		13 Oph,$\zeta$ Oph
152060			B2IVe		1.2
152478	6274		B3Vep+sh	0.78	0.97		2.18
152541			B7Ve		0.35
152979			B2IVe		0.1
153261	6304	245	B2IVe	0.95	0.99	$184\pm~5$	1.83	V828 Ara	Vega-type?
153295		588	B2IIIe		1
153879		246	B1.5Vne		0.8
154040		868	B2Ve		1.4	$175\pm18$
154111			B5Vne		0.4
154154			B2IVnne	0.80	1.7
154243		250	B3Vnne		2.0
154450		251	B0.5IVnne		1.7			V956 Sco
154911		871	B0.5IVne		0.35
155438			B6IVn		0.6
155806	6397	252	O8Ve	1.75	0.30	$116\pm23$			young?
155851		253	B1IIInne	$\approx$1.3	1.0
156325	6422	254	B5Vne	0.06	1.6			V1077 Sco
156468		255	B1Ve		0.65			V1078 Sco
157042	6451	258	B2IIIne	1.00	0.44	$318\pm28$	1.89	$\iota$ Ara
157832		259	B5Vnne			$277\pm18$	1.89	V750 Ara
158319		260	B4Ve		0.9	$234\pm~4$
158427	6510	261	B2Vne	1.18	0.58	$269\pm13$	2.23	$\alpha$ Ara
158643	6519		B9.5V	1.32	0.56	$210\pm31$	3.70	51 Oph	Vega-type
159489			B3IVep		0.35
160202		264	B7Ve		0.42
160886		885	B3Vnne		0.34	$328\pm32$		V2243 Oph
161103			B2IVe		0.95	$260\pm26$
161261			B9	$\approx$0.07	0.1	$297\pm30$		V2315 Oph
161543		887	B7.5e			$194\pm~4$		V2385 Oph
161711			B9III/IVe			$150\pm15$		GT Dra
161756	6621A		B6Ve	$\approx$0.4		$~85\pm30$		V3894 Sgr
161756	6621B		B8Ve			$154\pm15$		V3894 Sgr
161774		890	B5Vnne		0.55
161807			B3Vne		0.35	$287\pm29$	2.45
162352		892	B1.5Vne		0.5
162428		594	B7IV/Ve			$275\pm30$
162568		893	B3Vne		0.41
162732	6664		B6IVe+sh	0.4	0.76	$286\pm12$		V744 Her,88 Her
163454		276	B0.5Vpe		2.57
163848		899	B8Ve			$269\pm18$
163868		277	B3.5Vne		0.59	$297\pm29$		V3984 Sgr
164284	6712	278	B2IV/Ve+sh	0.77	0.52	$220\pm14$	2.10	66 Oph,V2048 Oph
164447	6720	279	B8Vne	0.43	0.45	$178\pm37$		V974 Her
164906		280	B1IV:pe:	1.25	0.55		4.98
165063			B5Vne		0.35
165285		281	B1.5Vne		1.5
165854			B9e			$242\pm10$
165921			B0Ve	0.41	0.48	$194\pm19$		V3903 Sgr
166014	6779		B9.5III		0.17	$142\pm18$	0.20	o Her,103 Her
166256		283	B9/A0e			$266\pm60$
166443		596	B1Ve		1.53	$246\pm25$		V4383 Sgr
166524			B0Vpe		1.09
166566		284	B1IIIe		0.51
166596	6804		B2.5IIIp			$207\pm~9$		V692 CrA
167128	6819		B3IIIep	0.41	0.24	$~66\pm~8$	1.30	QV Tel
168135		289	B8Ve			$250\pm20$
168797	6873	601	B3Ve	0.08	0.18	$251\pm~6$		NW Ser
168957		292	B4Ve		0.16	$~88\pm12$
169033	6881	602	B5Ve	$\approx$0.10	0.52	$210\pm~7$
169805		296	B2Ve		0.15	$274\pm~8$
170061		298	B0.5Vne		0.3
170235	6929	299	B2IVpe+sh	0.35	0.38	$163\pm16$	2.22	V4031 Sgr
170682			B4/5III			$121\pm12$		V3508 Sgr
170835			B2.5Vne		1.3

Appendix I. Observational data of classical Be stars

HD	HR	MWC	Sp	E(V-L)	$p_{\rm intr}$	$\overline{v\sin{i}}$	E(V-12)	other name	remarks
171348		302	B2Vne		0.27			V4400 Sgr
171406	6971		B4Ve	$\approx$0.30	0.28	$264\pm10$		V532 Lyr
171754		942	B4Ve			$138\pm~2$
171757		943	B2.5IIIne		0.60	$182\pm18$
171780	6984	604	B5Vne	-0.02	0.16	$256\pm~5$
172158		946	B7.5Ve		1.08	$232\pm26$
172252		947	B1Vnpe		0.35	$146\pm15$
172256		606	B1.5Vne		0.31			V4131 Sgr
173219		304	B1Vnep		1.83			V447 Sct
173370	7040		B9Ve	0.17	0.04	$261\pm~9$	0.56	4 Aql
173371		956	B9IIIe			$291\pm~4$
173637		607	B1IVe		1.95	$~98\pm10$
173948	7074	963	B2II/IIIe+sh		0.50	$177\pm12$		$\lambda$ Pav
174237	7084	608	B2.5Ve	0.96	0.18	$165\pm~9$	1.06	CX Dra
174513		609	B1Ve		1.8	$202\pm~5$		V457 Sct
174571		610	B3Vpe		1.2	$294\pm19$	1.92
174638	7106	306	B7Ve	0.22	0.45	$131\pm19$	1.14	$\beta$ Lyr
175863		308	B5Ve		0.8	$170\pm17$
175869	7158		B8IIIep	-0.02	0.46	$120\pm18$		64 Ser
176159		971	B9e			$218\pm~1$		V1437 Aql
177015		309	B5Vn		0.87	$236\pm24$
177291		973	B8Ve	1.44				V4409 Sgr	young?
177648		310	B2Ve		0.76
178175	7249	311	B2Ve	0.73	0.80	$156\pm26$	1.61	V4024 Sgr
178475	7262		B5V		0.21	$234\pm20$		$\iota$ Lyr,18 Lyr
179343		978	B9V+sh	0.15		$261\pm36$
179405			B3e			$234\pm~1$
181409	7335		B2IVe	$\approx$0	0.30
181615	7342	313	B2Vpe	3.03	0.68	$~69\pm~7$	5.66	$\nu$ Sgr	HAEBE?
182645	7378		B7III			$250\pm25$
183362	7403	318	B3Ve+sh	0.73	0.78	$220\pm10$	1.98	V558 Lyr
183656	7415	988	B6Ve+sh	0.25	0.54	$216\pm59$		V923 Aql
183914	7418	618	B8Ve	0.20		$220\pm25$		6 Cyg
184279		319	B0.5Ve+sh		0.9	$212\pm~9$		V1294 Aql
184915	7446		B0.5III	0.25	1.17	$229\pm~2$		$\kappa$ Aql
185037	7457	619	B8Vne	0.66	<0.05	$327\pm18$		11 Cyg
187350		621	B2IVne		0.85	$199\pm20$
187567	7554	322	B2.5IVe	$\approx$0.6	1.43		2.40	V1339 Aql
187811	7565	323	B2.5Ve+sh	0.58	0.15	$236\pm21$	0.67	12 Vul
187851			B2V		0.71	$234\pm23$		V396 Vul
189687	7647	624	B3IVe	0.61	0.28	$194\pm13$		25 Cyg,V1746 Cyg
190944		328	B1.5Vne	1.38	0.97	$162\pm16$	2.30
191531		328	B0.5IV			$~80\pm~3$
191610	7708	329	B3IVe	0.76	0.33	$290\pm31$	1.60	28 Cyg,V1624 Cyg
191639	7709		B1Ve	0.18	0.22	$235\pm24$		BE Cap
192044	7719	331	B7.5Ve	0.21	0.43	$287\pm~5$	0.89	20 Vul
192445		332	B2Ve		0.33	$242\pm14$
192685	7739		B3V	0.15	0.2	$206\pm19$	1.31	QR Vul
193009		336	B0Ve		2.1	$220\pm26$		V2113 Cyg
193182			B7IV/Ve+sh	0.25	0.15	$182\pm18$
193911	7789	341	B8IIIne+sh	0.11	0.43	$210\pm~6$	0.90	25 Vul
194244	7803		B9Ve	$\approx$0.1		$221\pm14$
194335	7807	343	B2IIIne	0.98	0.35	$328\pm23$		V2119 Cyg
194883		345	B2Ve		1.6			V2120 Cyg
195325	7836	1019	B9e+sh	$\approx$0	0.17	$250\pm36$		1 Del
195407		346	B0IV:pe	1.51	1.23	$259\pm26$
195554	7843	637	B9Vne	$\approx$0.05		$215\pm21$
196712	7890	350	B7IIIne	$\approx$0.10	0.05	$199\pm19$
197038			B7e			$194\pm~2$
197419	7927	1026	B2IV-Ve	$\approx$0.15	0.33	$115\pm~6$		V568 Cyg
197434		1027	B8e			$218\pm~9$
198183	7963	352	B5Ve	-0.02	0.32	$141\pm14$	-0.06	54 Cyg,$\lambda$ Cyg
198512		354	B1Vnnep		1.8	$178\pm18$		V2135 Cyg
198625	7983		B4Ve	0.12	0.25	$257\pm~6$	1.45	V2136 Cyg
198895		355	B1.5Ve		3.4	$172\pm18$	1.56	V417 Cep
198931		1032	B1Ve	1.60	1.92	$322\pm38$
199218	8009	356	B8Vnne	$\approx$0.25	0.24	$288\pm13$
199356		357	B2IVe	1.1	0.65	$354\pm24$		V2139 Cyg
200120	8047	359	B1.5Ve+sh	1.05	0.52	$281\pm33$	2.50	59 Cyg,V832 Cyg
200310	8053	360	B1Ve	0.14	0.51	$267\pm23$		60 Cyg,V1931 Cyg
201522			B0Ve		1.0
201733	8103	363	B4IVp	0.1	1.2	$270\pm27$		V2148 Cyg
202904	8146	364	B2.5Vne	1.09	0.46	$216\pm20$		66 Cyg
203025	8153	365	B2IIIe	$\approx$0.2	0.63	$130\pm28$
203064	8154		O7.5IIIe	0.50	0.23	$288\pm~6$		68 Cyg
203374		366	B0IVpe		1.1	$267\pm~3$	1.91
203467	8171	367	B2.5Ve+sh	0.73	0.03	$148\pm~2$	0.88	6 Cep,V382 Cep
203699			B3IV		0.18	$114\pm~1$
203731		369	B1Vne		1.0			V2153 Cyg
204116			B1Ve		1.18	$170\pm17$		V2155 Cyg
204185			B2e			$138\pm~3$
204722		370	B2Ve		1.95	$210\pm21$		V2162 Cyg
204860			B4Ve		0.2	$226\pm~8$		V2163 Cyg
205060		371	B6Ve		0.8	$214\pm12$
205551	8259	641	B9IIIe	$\approx$0		$166\pm27$
205618		642	B1.5Vne		0.95	$299\pm65$		V2166 Cyg
205637	8260	373	B2.5Ve+sh	0.81	1.20	$269\pm16$	0.54	$\epsilon$ Cap,39 Cap
206773		376	B0Vpe	1.6	0.29	$396\pm~6$	2.33
207232		377	B7V			$281\pm28$
208057	8356	644	B3Ve	0.04	0.24	$135\pm12$		16 Peg,OQ Peg
208220			B1IVe		0.48
208392		380	B1IIIe		0.34	$258\pm21$		EM Cep
208682	8375	381	B2.5Ve+sh	1.03	0.70	$269\pm19$	1.36
208886			B7III/IVe			$300\pm30$
209014	8386		B8Ve	0.30	0.15	$350\pm35$	0.67	$\eta$ PsA
209409	8402	384	B7IVe+sh	0.13	0.60	$275\pm32$	0.97	31 Aqr,o Aqr
209522	8408	650	B4IVne	-0.05	0.19	$265\pm27$		UU PsA
210129	8438	385	B7Vne	0.1	0.32	$201\pm32$		25 Peg
211835		652	B3Ve			$207\pm15$		V404 Lac
212044		386	B0Ve	$\approx$0.75	0.51	$162\pm~2$	1.87	V357 Lac
212076	8520	387	B1.5Vne	0.94	0.12	$121\pm13$		31 Peg,IN Peg
212571	8539	388	B1Ve+sh	0.34	0.94	$285\pm15$		52 Aqr,$\pi$ Aqr
212666		1059	B5.5e			$274\pm15$
212791		653	B6e			$170\pm~3$		V408 Lac
213088		389	B9e			$202\pm~7$

Appendix I. Observational data of classical Be stars

HD	HR	MWC	Sp	E(V-L)	$p_{\rm intr}$	$\overline{v\sin{i}}$	E(V-12)	other name	remarks
213129		654	B5.5e			$194\pm~1$
214167			B1.5Ve	$\approx$0.45	0.2	$295\pm30$		8 Lac A
214168	8603	390	B2Ve	1.19	0.18	$304\pm11$		8 Lac B
214748	8628	392	B8Ve	-0.03	0.34	$222\pm26$	0.63	18 PsA,$\epsilon$ PsA	Vega-type?
215227		656	B5ne		1.05	$262\pm26$
215605		658	B2IVnne		1.76
216057	8682	393	B5Vne	$\approx$-0.1	0.39	$283\pm24$
216200	8690		B3IV:e	$\approx$0.45	1.6	$199\pm17$		14 Lac,V360 Lac
216581			B2.5V			$282\pm28$
216851		660	B3Vn			$259\pm26$		V423 Lac
217050	8731	394	B3.5IIIpe+sh	0.72	0.77	$301\pm16$	2.21	EW Lac
217061			B1V	0.39	0.43	$226\pm23$
217543	8758	395	B3.5Vpe	$\approx$0.28	0.4	$297\pm17$
217675	8762		B6IIIpe+sh	0.07	0.13	$258\pm24$		o And,1 And
217891	8773	396	B5Ve	0.55	0.09	$118\pm14$	0.51	$\beta$ Psc,4 Psc
218393		397	B3IVne	1.0	0.33		1.25	KX And
218674			B3IV/Ve+sh	$\approx$0.1	0.36	$234\pm17$		KY And
219688	8858		B5V	-0.01	0.08	$271\pm11$		$\psi^{2}$ Aqr
220058		398	B1nep		0.1	$422\pm42$		V810 Cas
220116		399	B1e		1.77	$210\pm~4$		V811 Cas
223044			B3e			$146\pm~5$
223387		401	B0Ve		0.9	$202\pm~2$
224544	9068	406	B6IVe+sh	0	0.79	$155\pm39$
224559	9070	407	B3.5Vne	$\approx$0.15	0.24	$261\pm11$		LQ And
224686	9076		B8Ve	0.06	0.04	$261\pm14$	0.22	$\varepsilon$ Tuc	Vega-type?
224905			B1Vn		1.15	$258\pm26$
225095		409	B2IVne	$\approx$0.65	0.66	$152\pm10$
225985		995	B1Vpe	$\approx$1.2	0.9
227836		628	B0/2:V:e	$\approx$1.4	1.8$^\dagger$			V425 Cyg
228104		629	B1IVpe		0.95	$194\pm~3$
228438		333	B1Ve		0.15	$219\pm22$
228860			B0.5IVe		1.08
229171			B0.5IIIn		1.00
230579		975	B1.5IVne		1.35
232552		19	B0IVpe		1.09	$151\pm15$
235668		646	B2e			$247\pm28$		V2172 Cyg
235795		1057	B1:V:nne		1.2	$195\pm20$
236689			B1.5Vep	$\approx$0.35	0.25	$139\pm14$
236935		24	B1Vne		1.4	$159\pm16$
237056		720	B0.5Vpe		0.9	$160\pm17$
239758		643	B2Vne		1.1	$281\pm28$		V435 Cep
244894		761	B0.5Ve		1.45	$177\pm18$
245493		764	B2Vp		0.2	$206\pm21$
246878			B1Ve		1.1	$137\pm14$
248753		128	B1Vnne		1.7			V1167 Tau
249695		785	B1.5Vpnne		0.25
250028		786	B2Ve	$\approx$0.85	1.7
250163		517	B1Ve		1.8
250289			B2IIIe		1	$125\pm13$
254647		798	B0e		0.6
256577			B2IVp		1.3
276738		735	B6/7III/V			$220\pm22$
298298			B1Vpe		0.6
300584			B1Ve		1.6
302724			B2Ve		1
302838			B1Vne		0.4
305382			B2Ve		1.8
306209			B1Ve		0.6
306657			B2Ve			$224\pm22$
306791			B2IIIe			$154\pm15$
306793			B3Ve			$208\pm21$
306797			B5e			$288\pm29$		V855 Cen
306962			B1IIIne		1.2
306978			B2Ve			$144\pm14$
307350			B2IVne		0.5
312973			B0IVpe		0.9
316568			B2Vpe		0.8			AS 253
316587			B1Vne		1
316589			B2IVe		0.8
350559			B7IIIe			$186\pm19$
CD-30 5559			B3Vep		0.7
CPD-60 3087			B3Ve			$288\pm29$
CPD-60 3108			B2Vnpe			$208\pm21$
CPD-60 3122			B3nep			$224\pm22$		V845 Cen
CPD-60 3125			B2.5Ve			$232\pm23$
CPD-60 3126			B1.5Vne			$208\pm21$		V846 Cen
CPD-60 3128			B2.5Ve			$128\pm13$
CPD-60 3129			B2.5V			$184\pm18$
CPD-60 3144			B3Ve			$288\pm29$
CPD-60 4551			B1IIIne		0.71			V959 Cen
CPD-60 4708			B2Vne		0.38			V696 Cen
BD-13 2040		180	B7e		0.6
BD-13 4936		918	B1Vne		1.1
BD+35 1169		500	B1Vpe		0.9
BD+41 3731			B2e	0.17	0.85	$252\pm25$
BD+47 3302			B2Vnep		0.8			AS 458
BD+49 3735			B1.5Ve		0.4			AS 483,V397 Lac
BD+53 2964		1067	B2IVpnne		1
BD+55 589		446	B2IV/Ve			$295\pm30$
BD+56 473		441	B0.5IIIe	$\approx$0.55	1.5	$258\pm26$		V356 Per
BD+56 484		444	B1IIIe		0.5	$228\pm24$		V502 Per
BD+56 493			B1Vpe		0.8	$311\pm31$
BD+56 511			B3III	-0.08		$111\pm11$
BD+56 548		445	B1.5IIIe	1.78		$134\pm14$
BD+56 573		40	B2III/Ve	1.38		$335\pm34$
BD+56 579		710	B7IV/Ve			$311\pm31$
BD+56 582			B1III		1.7	$175\pm18$
BD+56 624		46	B3IIIe	$\approx$0.75	1.5	$234\pm23$
BD+57 515		28	B2ep		0.65	$218\pm22$
BD+57 607		55	Be		0.3	$298\pm30$
BD+58 458		42	B1pe		0.2	$418\pm42$
BD+59 2829		3	B0IVne		0.3
BD+61 2355			B7IV	0.36				BHJ 9
BD+61 2380			B9Ve	-0.1

$\dagger$ - observed polarization;

The columns in the Appendix I give the HD number or other identification, the MK spectral type of a star (as from the source catalogue), the calculated value of colour excess E(V-L) (in magnitudes), the calculated value of intrinsic V band polarization (in %), the weighted average value of $\overline{v\sin{i}}$ (in kms^-1), the average value of far IR excess (E(V-12)) at 12 $\mu$m, the common name or other designation of the star and any remarks. The references to the sources of the V and L photometry, the observed polarization and the individual values of $v\sin{i}$ are given below.

References to the sources of optical and IR photometry:
Dougherty et al. (1991); SIMBAD (1999); Gezari, D. Y. et al. (1999); Hoffleit & Jaschek (1982); Sterken (1990); Grillo et al. (1992); Hillenbrand et al. (1992);
References to the sources of optical polarimetry:
Poeckert et al. (1979); McLean & Brown (1978); Hall & Mikesell (1953); Ghosh et al. (1999); WUPPOL (1999); Leroy (1993); McDavid (1999); Serkowski et al. (1975); Quirrenbach (1997); Clarke & Bjorkman (1998); Huang et al. (1989); Bjorkman et al. (1997); Coyne et al. (1967); Jian & Bhatt (1995); Yudin & Evans (1998); Vosnchinnikov & Marchenko (1982); Bjorkman & Schulte-Ladbeck (1994); McDavid (2000); Heiles (2000).
References to the sources of $v\sin{i}$ data:
SIMBAD (1999); Balona (1990); Slettebak (1982); van den Ancker et al. (1998); Corporon & Lagrange (1999); Elias et al. (1997); Bjorkman et al. (1997); WUPPOL (1999); Brown & Verschueren (1997); Howarth et al. (1997); Prinja (1993); Hirata (1993); Quirrenbach (1993); Mennickent et al. (1994); Ballereau et al. (1995); Hoffleit & Jaschek (1982); Halbedel (1996); Uesugi & Fukuda (1982); Kaper (1999); Henrichs (1999); Dachs et al. (1981); Steele et al. (1999); Carpenter et al. (1984); Levato & Maloroda (1970); Dachs et al. (1981); Coté & Waters (1987); Lyubimkov et al. (1997); Balona (1995); Short & Bolton (1994); Rivinius et al. (1999); Prosser (1992); Snow (1981); Snow (1981); Ghosh (1988); Slettebak et al. (1997); Briot (1986); Chosh et al. (1999); Doazan et al. (1981); Wolff et al. (1978); Wolff et al. (1982); Finkenzeller (1985); Balona et al. (1992); Turner et al. (1978); Slettebak (1966); Jaschek & Jaschek (1993); Conti & Ebberts (1977); Herrero (1994); Penny (1996); Bernacca & Perinotto (1970); Clark & Steele (2000)

Appendix II. Observed, interstellar and intrinsic polarization

		observed		interstellar		intrinsic				observed		interstellar		intrinsic
		polarization		polarization		polarization				polarization		polarization		polarization
HD	D, pc	$p_{V}\pm\sigma_{\rm p}$, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$	HD	D, pc	$p_{V}\pm\sigma_{\rm p}$, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$
144	301	$0.62\pm0.05$	73	0.73	85	0.30	24	47359	1820	$1.01\pm0.20$	8	0.5	170	0.7	20
4180	278	$0.70\pm0.13$	85	0.5	90	0.23	74	47761	751	$2.30\pm0.20$	169	1.8	175	0.6	152
5394	188	$0.60\pm0.25$	111	0.3	90	0.43	125	48282	425	$0.83\pm0.04$	165	0.6	35	1.0	146
6811	226	$0.74\pm0.12$	90	0.6	90	0.14	90	48917	800					0.51$\otimes$	18
7636	559	$2.59\pm0.20$	80	2.3	85	0.51	55	49330	298	$1.71\pm0.18$	133	0.8	150	1.14	121
10144	44	$0.04\pm0.00$	136	0		0		49336	444	$1.30\pm0.10$	2	0.9	50	1.65	166
10516	220	$1.80\pm0.30$	28	0.4	105	2.17	26	49787	1445	$0.00\pm0.20$		0.7	160	0.7	70
11415	136	$0.40\pm0.17$	125	0.2	105	0.28	139	49888	502	$0.00\pm0.20$		0.50	130	0.50	40
12302	340	$2.34\pm0.11$	90	1.54	90	0.80	89	49977	676	$2.17\pm0.20$	152	0.30	130	1.95	155
12856	1700	$1.20\pm0.20$	108	2	100	0.91	179	50013	242	$0.31\pm0.10$	106	0		0.3	106
13051	2041	2.21	88	2.2	100	0.92	49	50083	500	$1.24\pm0.20$	24	1	170	1.25	47
13758	1620	$4.61\pm0.18$	108	3.5	115	1.48	91	50658	253	$0.62\pm0.12$	16	0.2	50	0.6	4
13890	2000	$3.46\pm0.18$	107	3.3	115	0.95	71	50696	1740	$0.23\pm0.18$	139	0.8	150	0.6	64
14422	1950	$3.09\pm0.18$	117	3.5	115	0.47	11	51354	398	$0.92\pm0.20$	140	0.2	15	1.0	135
14605	1820	$3.78\pm0.20$	118	3.5	115	0.47	144	51480	160	$0.37\pm0.20$	173	0.4	130	0.55	17
15450	725	$3.87\pm0.20$	122	3.5	115	0.97	152	52244	1660	$0.69\pm0.18$	13	0.7	120	1.33	22
18552	341	$\odot$				0.38	104	52437	407	$0.16\pm0.10$	81	0.2	100	0.12	136
18877	1820	$2.12\pm0.20$	126	2	120	0.45	160	52721	909	$1.14\pm0.15$	17	0.4	135	1.4	24
19243	617	$2.21\pm0.20$	129	1.9	119	0.78	158	52812	460	$0.88\pm0.20$	39	0.3	160	1.05	46
20017	503	$2.86\pm0.20$	126	2.6	120	0.62	156	52918	342	$0.00\pm0.20$		0.4	130	0.4	40
20134	394	$1.75\pm0.20$	110	2.4	120	0.90	48	53367	830	$0.60\pm0.20$	30	0.2	130	0.75	32
20336	246	$0.43\pm0.08$	142	1.0	128	0.65	29	54309	380	$0.65\pm0.10$	1	0.3	30	0.55	167
20418	147	$0.34\pm0.00$	101	0.25	110	0.13	83	54464	1000	$0.00\pm0.20$		0		0
21212	407	$4.06\pm0.20$	136	2.7	123	2	154	54786	2100	$0.78\pm0.18$	116	0.5	170	1	102
21362	170	$0.31\pm0.00$	112	0.4	110	0.1	14	55135	420	$0.00\pm0.20$		0.3	130	0.30	40
21455	177	$0.55\pm0.32$	106	0.5	111	0.10	78	55538	1800	$1.15\pm0.20$	171	0.5	170	0.65	172
21551	267	$0.51\pm0.15$	127	1.3	120	0.85	26	55606	2500	$0.55\pm0.18$	60	0		0.55	60
21641	220					0.26$\otimes$	129	56014	483	$0.62\pm0.05$	85	0.4	90	0.25	76
21650	314	$0.65\pm0.20$	151	0.9	130	0.6	17	56139	283	$0.20\pm0.15$	60	0.2	80	0.14	25
22192	215	$0.80\pm0.15$	45	1.0	117	1.7	35	56806		$0.65\pm0.20$	178	0.7	165	0.3	41
22298	360	$2.67\pm0.12$	119	1.55	120	1.12	118	57150	260	$0.70\pm0.04$	87	0.15	30	0.77	92
22780	248	$0.51\pm0.05$	144	0.6	110	0.47	174	57219	256	$0.12\pm0.04$	32	0.15	30	0.03	113
23016	125	$0.60\pm0.04$	87	0.3	120	0.55	72	57386	1250	$0.55\pm0.18$	141	0.4	135	0.18	154
23302	114	$0.32\pm0.05$	128	0.3	120	0.09	162	57393	2400	$0.49\pm0.15$	78	0.3	130	0.6	65
23478	239	$1.71\pm0.20$	29	0.8	7	1.3	42	57539	483	$0.00\pm0.20$		0.5	120	0.5	30
23480	110	$0.37\pm0.02$	135	0.3	120	0.18	162	58011	1085	$0.36\pm0.10$	111	0.35	120	0.1	73
23552	182	$0.55\pm0.05$	140	0.4	140	0.16	154	58050	691	$0.06\pm0.00$	78	1.1	25	1.12	156
23630	113	$0.18\pm0.02$	117	0.3	120	0.12	34	58343	290	$0.88\pm0.20$	164	0.5	170	0.40	157
23800	427	$2.67\pm0.20$	122	1.55	130	1.26	112	58715	52	$0.14\pm0.13$	90	0		0.14	90
23862	119	$0.43\pm0.03$	90	0.3	120	0.38	69	58978	435	$1.10\pm0.30$	161	0.25	160	0.85	161
24479	100	$0.06\pm0.03$	146	0.27	136	0.21	43	59094	870	$0.51\pm0.18$	118	0.6	110	0.18	173
24534	826	0.88	43	1	57	0.48	178	59498	1050	$0.64\pm0.10$	115	0.6	120	0.12	82
25348	1260	$1.34\pm0.20$	124	0.5	95	1.16	135	60848	498	$0.16\pm0.25$	56	0.25	25	0.22	96
25799	360	$0.37\pm0.20$	102	0.8	95	0.45	179	60855	508					0.31$\oslash$
25940	170	$0.25\pm0.10$	145	0.25	165	0.17	110	61355		$0.41\pm0.10$	102	0.8	140	0.8	65
26356	210	$0.00\pm0.20$		0.06	136	0.06	46	62532	1580	$0.46\pm0.18$	127	0.45	165	0.55	101
26420	182	$0.78\pm0.20$	137	0.3	165	0.66	126	62654	1380	$1.63\pm0.10$	163	1	170	0.7	153
26906	550	$2.17\pm0.20$	144	1.1	155	1.22	134	62753	298	$0.72\pm0.10$	15	0.15	45	0.65	9
28497	483	$0.90\pm0.20$	128	0.3	180	1.02	120	62780	4170	$1.65\pm0.10$	128	0.5	70	1.90	135
29866	166	$1.98\pm0.20$	6	0.8	10	1.19	3	63150	1820	$2.01\pm0.10$	137	1	160	1.5	123
30076	412	$0.40\pm0.20$	169	0.3	180	0.17	148	63359	2630	$0.26\pm0.06$	25	0.8	150	0.9	52
32190	575	$0.78\pm0.20$	80	0.7	80	0.08	80	63462	315	$0.00\pm0.20$		0.2	100	0.2	10
32343	206	$0.51\pm0.05$	168	0.4	165	0.12	178	64298	524	$0.75\pm0.10$	122	0.1	100	0.68	125
32990	302	$1.34\pm0.20$	88	0.7	80	0.69	96	65719	1660	$0.65\pm0.20$	14	0.5	15	0.15	11
32991	316	$0.80\pm0.35$	95	0.7	80	0.40	126	65875	315	$0.15\pm0.20$	15	0		0.15	15
33152	955	$1.57\pm0.20$	4	1.4	150	1.67	30	66194	315	$0.65\pm0.05$	138	0.35	145	0.32	130
33232	324	$1.15\pm0.20$	140	1.0	165	0.92	112	69168	310	$0.26\pm0.10$	80	0.1	80	0.16	80
33328	538	$0.11\pm0.00$	78	0.3	110	0.24	31	69404	365	$0.56\pm0.10$	120	0.2	100	0.4	129
33461	372	$1.84\pm0.20$	165	1.2	165	0.64	165	69425	2290	$0.72\pm0.10$	131	1.5	110	1.11	7
33599	794	$0.68\pm0.05$	30	0.1	90	0.7	30	71934	830	$0.82\pm0.10$	13	0.25	180	0.6	18
33604	603	$1.29\pm0.20$	168	1.5	160	0.44	43	72014	315	$0.56\pm0.10$	126	0.1	130	0.46	125
33988	240	$0.78\pm0.20$	161	1	160	0.22	66	72063	2000	$1.17\pm0.10$	107	0.5	130	0.9	95
34921	1515	$3.73\pm0.20$	146	2	160	2	133	72067	490	$0.02\pm0.02$	160	0.2	180	0.2	92
34959	382	$0.94\pm0.03$	65	0.3	80	0.7	59	74401	1510	$0.74\pm0.10$	81	0.6	160	1.3	76
35165	581	$0.77\pm0.04$	6	0.1	10	0.7	5	75658	690	$0.65\pm0.10$	138	0.45	160	0.45	20
35345	1115	$2.95\pm0.20$	155	2.5	160	0.65	134	76534	411	$0.49\pm0.05$	124	1	130	0.53	46
35411	276	$0.10\pm0.06$	168	0.1	70	0.20	165	77032	1200	$1.74\pm0.10$	123	0.6	120	1.14	124
35439	340	$0.35\pm0.07$	122	0.3	80	0.46	142	77320	307	$0.85\pm0.10$	130	0.57	70	1.24	142
36012	391	$0.24\pm0.03$	72	0.3	80	0.10	12	78764	251	$0.28\pm0.04$	125	0.2	125	0.08	125
36576	575	$0.11\pm0.00$	75	1.15	25	1.17	113	79778	830	$1.00\pm0.10$	145	0.8	145	0.2	145
37115	543	$0.37\pm0.03$	89	0.4	60	0.37	121	80834	795	$1.55\pm0.10$	111	1.0	100	0.73	127
37202	128	$1.40\pm0.15$	31	0.3	20	1.13	34	81654	724	$0.47\pm0.10$	131	1.1	120	0.7	23
37330	493	$0.66\pm0.03$	63	0.4	65	0.26	60	82830	3980	$0.86\pm0.10$	128	1.4	10	2.0	110
37490	498	$0.43\pm0.10$	48	0.45	65	0.27	9	83043	1050	$1.16\pm0.13$	162	2.0	155	0.92	56
37657	363	$1.66\pm0.20$	174	0.4	160	1.32	178	83060	1260	$1.21\pm0.10$	108	1.7	110	0.5	25
37795	82	$0.15\pm0.05$	109	0		0.15	109	83597	2400	$1.24\pm0.07$	151	3.0	155	1.78	68
37967	337	$0.78\pm0.20$	8	0.3	85	1.1	4	83953	152	$0.27\pm0.03$	176	0.05	75	0.32	174
38010	262	$1.77\pm0.08$	158	1.5	5	0.33	131	84523	525	$0.26\pm0.10$	122	0.5	120	0.25	28
38087	199	$2.53\pm0.06$	118	0.2	65	2.5	122	84567	550	$0.75\pm0.04$	74	0.5	75	0.25	73
38191	2400	$1.89\pm0.20$	156	1.5	150	0.5	174	86272	1050	$0.89\pm0.10$	118	0.8	100	0.5	149
38708	724	$1.15\pm0.18$	167	2.2	170	1.06	83	86612	193	$0.04\pm0.04$	10	0.03	10	0.01	10
39340	794	$2.30\pm0.20$	171	2.2	170	0.13	10	86689	3800	$0.20\pm0.10$	168	0.6	110	0.7	13
39478	912	$0.00\pm0.12$		2	170	2.0	80	87380	2890	$1.34\pm0.10$	127	1.0	130	0.2	84
40193	650	$0.48\pm0.07$	70	0.4	40	0.4	96	87543	250	$0.50\pm0.10$	107	0.3	105	0.20	110
40978	460	$1.06\pm0.20$	6	0.5	180	0.58	11	87901	83.5	$0.06\pm0.12$	46	0.1	80	0.1	23
41335	282	$0.55\pm0.28$	147	0.05	60	0.60	147	89080		$0.20\pm0.05$	100	0.6	120	0.45	38
42054	323	$0.01\pm0.05$	35	0		0		90187	3020	$1.84\pm0.10$	115	1.5	120	0.45	96
42545	238	$0.00\pm0.20$		0.1	45	0.1	135	90563	1100	$1.77\pm0.10$	135	1.5	120	0.9	164
43285	228	$0.03\pm0.07$	61	0.1	45	0.08	129	91120	149	$0.23\pm0.05$	89	0		0.23	89
43544	256	$0.14\pm0.06$	46	0.2	45	0.05	133	91188	794	$0.34\pm0.10$	105	1.2	150	1.25	68
43703	1400	$1.75\pm0.20$	142	0.5	155	1.3	137	91465	152	$0.90\pm0.10$	74	0.35	140	1.16	68
44458	513	$0.84\pm0.08$	79	0.35	50	0.72	91	92027	870	$1.15\pm0.10$	111	0.5	100	0.71	119
44506	617					0.39$\otimes$	92	92061	2890	$0.94\pm0.10$	99	0.9	100	0.05	80
44637	398	$1.84\pm0.20$	1	1.3	175	0.6	14	92714	1200	$0.37\pm0.10$	31	0.8	100	1.1	16
45314	457	$1.38\pm0.20$	171	1.4	175	0.2	124	92938	140	$0.17\pm0.04$	147	0.4	140	0.24	45
45542	154	$0.22\pm0.12$	17	0.15	5	0.10	35	93030	135	$0.38\pm0.10$	139	0.4	140	0.02	67
45626	457	$2.03\pm0.20$	157	2	175	1.25	122	93237	307	$0.74\pm0.04$	137	0.5	140	0.25	131
45725	212	$0.39\pm0.15$	0.8	0.2	165	0.24	14	93561		$0.89\pm0.10$	105	0.8	140	1.0	80
45726	212	$0.08\pm0.05$	155	0.2	165	0.13	81	93683	1260	$0.98\pm0.10$	146	1	80	1.8	158
45910	760	1.06	150	0.5	170	0.70	135	94648	1260	$1.29\pm0.10$	78	1	80	0.3	72
45995	435	$0.78\pm0.25$	132	0.4	180	0.91	119	96357	830	$1.50\pm0.10$	130	1	100	1.3	150
46380	575	$1.80\pm0.20$	178	0.6	10	1.28	172	96864	1380	$1.34\pm0.10$	150	1.5	120	1.4	3
47054	256					0.34$\otimes$	71	98624	2290	$0.79\pm0.10$	119	1	80	1.1	149

Appendix II. Observed, interstellar and intrinsic polarization

		observed		interstellar		intrinsic				observed		interstellar		intrinsic
		polarization		polarization		polarization				polarization		polarization		polarization
HD	D, pc	$p_{V}\pm\sigma_{\rm p}$, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$	HD	D, pc	$p_{V}\pm\sigma_{\rm p}$, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$
98927	1510	$1.48\pm0.10$	63	1.0	80	0.8	43	160886	1000	$0.97\pm0.20$	59	1.3	57	0.34	141
99354	2750	$1.39\pm0.10$	72	1.3	90	0.84	35	161103	830	$4.83\pm0.12$	172	4.0	175	0.95	158
100199	2290	$0.86\pm0.00$	84	1.3	90	0.5	11	161261	452	$0.81\pm0.05$	82	0.75	85	0.1	57
100324	912	$4.26\pm0.10$	136	3.5	120	2	164	161774		$1.44\pm0.10$	164	0.9	160	0.55	171
100889	93	$0.03\pm0.03$	33	0.07	35	0.04	128	161807	383	$1.73\pm0.10$	177	1.4	175	0.35	5
102742	1200	$1.10\pm0.10$	86	1.0	90	0.15	61	162352		$2.01\pm0.10$	173	1.5	175	0.5	166
105435	121	$0.33\pm0.07$	137	0		0.33	137	162568	253	$1.06\pm0.10$	179	0.65	177	0.41	3
105521	330	$0.31\pm0.04$	30	0.07	111	0.38	28	162732	368	$0.31\pm0.05$	67	0.45	160	0.76	69
105675	2890	$1.74\pm0.10$	59	1.2	90	1.55	38	163454	1200	$1.67\pm0.15$	50	1.3	170	2.57	63
105753	3980	$0.30\pm0.10$	74	1.2	90	0.7	41	163868	600	$0.45\pm0.10$	141	1.0	150	0.59	67
107348	118	$0.10\pm0.05$	132	0.07	140	0.04	22	164284	207	$0.85\pm0.25$	92	0.5	75	0.52	108
109387	153	$0.25\pm0.10$	24	0.1	60	0.24	12	164447	498	$0.29\pm0.04$	146	0.55	175	0.45	101
110335	317	$0.46\pm0.04$	89	1.45	75	1.05	160	164906	1050	$0.37\pm0.20$	174	0.9	180	0.55	94
110432	301	$1.77\pm0.11$	81	0.45	80	1.3	81	165063	377	$0.70\pm0.10$	147	0.8	160	0.35	100
110863	760	$4.47\pm0.06$	93	2.8	85	1.94	105	165285	1580	$1.04\pm0.14$	56	0.6	120	1.5	47
112078	122	$0.05\pm0.04$	169	0.05	125	0.07	13	165921	541	$0.74\pm0.04$	100	0.5	120	0.48	79
112091	111	$0.17\pm0.04$	128	0.05	130	0.12	127	166014	106	$0.17\pm0.01$	179	0		0.17	179
112999	336	$1.82\pm0.10$	78	2.2	90	0.9	27	166443	1510	$1.63\pm0.20$	145	0.5	110	1.53	154
113120	483	$1.18\pm0.15$	115	1.1	105	0.4	149	166524	2290	$0.51\pm0.18$	105	0.8	50	1.09	127
113605	1445	$1.71\pm0.10$	80	2.0	90	0.7	28	166566	1380	$0.97\pm0.20$	45	1.0	30	0.51	84
113902	99	$0.03\pm0.04$	107	0.05	130	0.04	58	167128	228	$0.98\pm0.10$	3	0.9	10	0.24	151
114200	3630	$1.40\pm0.10$	92	1.2	100	0.4	64	168797	360	$0.78\pm0.20$	64	0.60	65	0.18	61
114800	382	$0.72\pm0.10$	109	1.8	75	1.7	153	168957	513	$0.78\pm0.20$	176	0.65	180	0.16	159
114981	935	$0.14\pm0.04$	32	0.15	35	0.02	152	169033	188	$1.06\pm0.20$	111	0.8	125	0.52	88
116781	1995	$1.46\pm0.10$	81	2.2	75	0.8	155	169805	1260	$0.83\pm0.10$	141	1	140	0.15	45
117111	1050	$0.89\pm0.15$	78	1.5	80	0.6	52	170061		$1.19\pm0.19$	71	1.5	70	0.3	156
117357	1995	$3.03\pm0.10$	74	2.2	75	0.8	71	170235	603	$0.82\pm0.10$	175	0.5	4	0.38	163
118246	660	$0.29\pm0.05$	83	0.25	80	0.05	99	170835	427	$1.82\pm0.10$	21	0.9	180	1.3	35
120324	162	$0.07\pm0.05$	169	0.07	170	0		171348	476	$1.08\pm0.09$	174	0.9	180	0.27	152
120991	630	$0.47\pm0.10$	84	0.4	80	0.09	102	171406	455	$0.28\pm0.20$	13	0.5	180	0.28	19
121556	760	$0.47\pm0.47$	129	1.6	80	1.73	162	171757	3980	$0.43\pm0.10$	36	0.8	60	0.6	166
121847	104	$0.30\pm0.05$	65	0.25	60	0.04	9	171780	318	$0.28\pm0.20$	171	0.4	180	0.16	106
122450	833	$1.37\pm0.15$	129	1.7	80	0.66	163	172158	2090	$0.54\pm0.10$	143	1.1	180	1.08	104
122669	1510	$2.35\pm0.10$	64	2.2	70	0.48	31	172252	725	$4.65\pm0.20$	148	4.5	150	0.35	117
122691	1660	$1.82\pm0.10$	87	2.2	70	1.2	132	172256	760	$0.81\pm0.12$	169	0.7	180	0.3	140
124367	149	$0.11\pm0.08$	120	0.3	70	0.35	150	173219	1260	$0.65\pm0.20$	41	2.0	45	1.83	156
124771	169	$0.67\pm0.04$	112	0.3	110	0.37	114	173370	145	$0.13\pm0.00$	81	0.12	73	0.04	118
125159		$1.12\pm0.10$	62	1	70	0.3	33	173637	1900	$1.15\pm0.20$	87	1.7	45	1.95	117
127489	535	$1.58\pm0.10$	62	2	65	0.46	166	173948	556	$0.62\pm0.10$	21	0.15	20	0.50	20
127972	94.6	$0.37\pm0.07$	174	0.1	75	0.47	172	174237	667	$0.18\pm0.15$	80	0		0.18	80
129954	351	$1.06\pm0.04$	65	1.3	80	0.67	17	174513	2180	$0.00\pm0.20$		1.8	45	1.8	135
131492	256	$0.46\pm0.04$	65	0.9	75	0.5	175	174571	460	$2.90\pm0.20$	80	1.7	80	1.2	80
133738	115	$0.95\pm0.10$	52	0.15	60	0.81	51	174638	270	$0.60\pm0.15$	155	0.15	155	0.45	155
134401	1260	$0.62\pm0.10$	13	1.1	55	1.2	161	175863	400	$0.00\pm0.20$		0.8	110	0.8	20
135160	558	$1.18\pm0.04$	54	1.5	60	0.4	168	175869	338	$0.04\pm0.04$	48	0.5	60	0.46	151
136415	156	$0.56\pm0.10$	83	0.3	65	0.36	97	177015	800	$0.51\pm0.20$	0	0.8	40	1.87	148
136556	1900	$2.00\pm0.10$	34	2	55	1.45	179	177648	279	$1.15\pm0.20$	40	0.5	25	0.76	50
137387	313	$0.93\pm0.04$	122	0.65	135	0.45	102	178175	469	$0.35\pm0.30$	0	0.72	45	0.80	148
137432	128	$0.09\pm0.04$	91	0.1	65	0.08	125	178475	255	$0.51\pm0.20$	138	0.3	140	0.21	135
137518	4170	$0.65\pm0.10$	41	1.5	55	1.0	154	181409	568	$0.37\pm0.20$	163	0.5	145	0.30	31
138749	95.3	$0.20\pm0.20$	90	0.05	100	0.15	87	181615	513	$1.38\pm0.20$	172	0.7	172	0.68	172
138769	133	$0.08\pm0.04$	66	0.05	45	0.05	85	183362	525	$0.75\pm0.20$	42	0.3	0	0.78	53
139790	1445	$0.80\pm0.10$	36	2.5	50	1.8	146	183656	295	$0.90\pm0.22$	60	0.4	70	0.54	53
140336	3800	$1.17\pm0.10$	49	1.7	55	0.6	156	184279	600	$0.00\pm0.20$		0.9	100	0.9	10
140784	119	$0.04\pm0.04$	32	0.1	70	0.1	171	184915	446	$1.35\pm0.33$	172	2.4	180	1.17	99
140926	578	$1.64\pm0.10$	43	1.5	50	0.42	11	185037	226	$0.15\pm$	77	0.15	75	<0.05	121
140946	1095	$1.31\pm0.10$	60	1.5	50	0.5	111	187350	322	$0.78\pm0.20$	82	1.6	90	0.85	7
141637	160	$0.97\pm0.20$	34	0.4	45	0.62	27	187567	476	$0.51\pm0.20$	114	1.35	70	1.43	31
141926	1445	$1.74\pm0.10$	66	2	50	1.05	110	187811	190	$0.00\pm0.20$		0.15	5	0.15	96
142184	120	$0.65\pm0.20$	64	0.5	60	0.17	76	187851	368	$1.11\pm0.18$	24	0.55	40	0.71	12
142237	1095	$1.41\pm0.10$	55	2	50	0.65	129	189687	483	$0.20\pm0.10$	110	0.4	90	0.28	166
142926	148	$0.16\pm0.05$	160	0		0.16	160	190944	1510	$1.34\pm0.20$	24	0.4	15	0.97	28
142983	157	$0.85\pm0.15$	138	0.65	110	0.73	162	191610	264	$0.35\pm0.15$	56	0.15	90	0.33	43
144320	1260	$1.15\pm0.10$	48	2	50	0.85	143	191639	860	$0.51\pm0.20$	4	0.3	180	0.22	10
144970	1660	$2.86\pm0.13$	44	2.5	50	1.0	32	192044	350	$0.15\pm0.05$	160	0.30	90	0.43	173
145846	630	$3.24\pm0.11$	55	3	50	0.85	86	192445	1860	$0.37\pm0.20$	174	0.7	175	0.33	86
146444	760	$1.46\pm0.10$	24	2.5	35	1.25	138	192685	368	$0.00\pm0.20$		0.2	57	0.2	148
146463	725	$0.68\pm0.10$	31	2	50	1.5	148	193009	860	$0.65\pm0.20$	75	1.5	170	2.1	78
147274	2290	$1.51\pm0.10$	35	1	45	0.65	20	193182	285	$0.47\pm0.10$	134	0.35	140	0.15	119
147302	1320	$2.32\pm0.10$	52	2	50	0.35	60	193911	552	$0.00\pm0.05$		0.43	57	0.43	147
147756	2090	$1.39\pm0.10$	24	1.5	25	0.1	127	194335	274	$0.00\pm0.20$		0.35	70	0.35	160
148184	150	$0.40\pm0.05$	110	0.7	125	0.41	50	194883	769	$0.00\pm0.20$		1.6	160	1.6	70
148259	457	$1.27\pm0.10$	19	1.6	20	0.3	114	195325	174	$0.36\pm0.07$	63	0.2	70	0.17	55
148877		$0.83\pm0.10$	39	1	45	0.25	156	195407	1360	$1.80\pm0.20$	44	1.6	65	1.23	14
149298	1200	$0.82\pm0.10$	156	0.8	35	1.3	141	196712	377	$0.35\pm0.04$	95	0.4	95	0.05	2
149313	366	$1.69\pm0.10$	30	1.4	25	0.39	49	197419	360	$0.00\pm0.20$		0.33	45	0.33	135
149671	211	$0.57\pm0.04$	48	0.5	40	0.16	89	198183	360	$0.17\pm0.12$	81	0.33	45	0.32	120
149729	1380	$1.24\pm0.10$	49	0.9	45	0.38	60	198512	1000	$1.66\pm0.20$	126	2.4	150	1.8	82
149757	140	$0.70\pm0.60$	126	0.7	90	0.82	153	198625	282	$0.00\pm0.20$		0.25	65	0.25	155
152060	2190	$0.73\pm0.10$	23	0.8	35	1.2	175	198895	397	$5.39\pm0.20$	149	2.5	135	3.4	159
152478	230	$1.26\pm0.12$	19	0.3	25	0.97	17	198931	752	$4.00\pm0.29$	147	2.8	160	1.92	143
152541	380	$0.93\pm0.10$	21	0.6	26	0.35	14	199218	207	$0.37\pm0.20$	60	0.3	80	0.24	43
152979	1660	$1.69\pm0.10$	26	1.6	25	0.1	37	199356	645	$2.17\pm0.20$	151	2	160	0.65	118
153261	363	$1.57\pm0.05$	24	0.8	40	0.99	11	200120	345	$0.45\pm0.10$	22	0.3	65	0.52	5
153295	2290	$1.84\pm0.10$	33	0.8	35	1	32	200310	418	$0.45\pm0.05$	100	0.33	60	0.51	120
153879	1660	$1.62\pm0.10$	13	1.0	25	0.8	178	201522	418	$0.97\pm0.20$	15	0.33	60	1.0	6
154040	630	$1.15\pm0.50$	86	0.6	35	1.4	98	201733	336	$0.97\pm0.20$	172	0.3	62	1.2	168
154111	1380	$1.29\pm0.10$	35	1.6	30	0.4	104	202904	276	$0.33\pm0.05$	111	0.15	40	0.46	117
154154	1820	$2.72\pm0.10$	21	1.0	25	1.7	18	203025	525	$0.18\pm0.20$	11	0.8	20	0.63	113
154243	1050	$1.41\pm0.10$	22	0.8	90	2.0	14	203064	1030	$0.37\pm0.20$	55	0.15	65	0.23	49
154450	1820	$1.00\pm0.20$	90	1.25	140	1.7	67	203374	830	$0.37\pm0.20$	166	1	35	1.1	134
154911	2190	$1.66\pm0.15$	149	1.65	155	0.35	107	203467	331	$0.50\pm0.25$	122	0.5	120	0.03	166
155438	2750	$1.68\pm0.10$	28	1.1	25	0.6	33	203699	600	$0.18\pm0.20$	33	0		0.18	33
155806	660	$0.84\pm0.09$	150	0.7	140	0.30	165	203731	309	$0.74\pm0.20$	167	0.27	65	1.0	164
155851	1445	$2.18\pm0.21$	146	1.25	140	1.0	154	204116	725	$2.95\pm0.20$	22	1.8	25	1.18	17
156325	386	$2.01\pm0.08$	156	0.5	140	1.6	161	204722	780	$1.20\pm0.20$	10	2.0	45	1.95	153
156468	1000	$1.00\pm0.10$	9	1.5	180	0.65	75	204860	483	$0.00\pm0.20$		0.2	45	0.2	135
157042	221	$0.26\pm0.04$	151	0.25	30	0.44	136	205060	461	$0.97\pm0.20$	32	0.22	50	0.8	27
158319	725	$0.78\pm0.20$	70	1.5	55	0.9	132	205618	990	$1.24\pm0.20$	26	0.6	50	0.95	12
158427	74	$0.61\pm0.04$	174	0.15	30	0.58	166	205637	203	$1.20\pm0.25$	156	0		1.20	156
158643	131	$0.46\pm0.10$	30	0.1	120	0.56	30	206773	498	$2.03\pm0.20$	169	1.75	170	0.29	163
159489	3800	$1.32\pm0.10$	22	1.0	20	0.35	27	208057	157	$0.09\pm0.04$	119	0.19	32	0.24	144
160202	532	$1.32\pm0.13$	161	1.7	165	0.42	87	208220	2600	$1.94\pm0.20$	42	2.0	35	0.48	87

Appendix II. Observed, interstellar and intrinsic polarization

		observed		interstellar		intrinsic
		polarization		polarization		polarization
HD	D, pc	$p_{V}\pm\sigma_{\pm p}$, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$	pV, %	$\theta$, $^\circ$
208392	910	$1.34\pm0.20$	44	1	45	0.34	41
208682	369	$1.00\pm0.30$	140	0.5	120	0.70	154
209014	311	$0.15\pm0.10$	15	0		0.15	15
209409	117	$0.52\pm0.05$	177	0.2	125	0.60	6
209522	346	$0.19\pm0.04$	19	0		0.19	19
210129	214	$0.28\pm0.04$	7	0.12	55	0.32	176
212044	860	$1.26\pm0.20$	52	1.7	47	0.51	124
212076	298	$0.25\pm0.10$	98	0.26	112	0.12	58
212571	338	$0.93\pm0.75$	168	0.3	127	0.94	177
214167	196	$0.00\pm0.20$		0.20	110	0.2	20
214168	77	$0.13\pm0.20$	30	0.05	120	0.18	30
214748	228	$0.09\pm0.10$	155	0.30	100	0.34	3
215227	256	$0.66\pm0.12$	28	0.60	83	1.05	12
215605	1995	$1.45\pm0.18$	40	2.6	60	1.76	166
216057	185	$0.38\pm0.03$	44	0.75	50	0.39	146
216200	333	$0.89\pm0.20$	145	0.8	75	1.6	154
217050	337	$1.55\pm0.20$	75	0.8	75	0.77	67
217061	795	$1.22\pm0.20$	88	1.5	95	0.43	27
217543	273	$0.00\pm0.20$		0.4	85	0.4	175
217675	212	$0.34\pm0.12$	90	0.45	85	0.13	161
217891	151	$0.28\pm0.06$	109	0.22	102	0.09	128
218393	446	$0.72\pm0.10$	55	0.9	65	0.33	179
218674	525	$0.65\pm0.05$	90	0.9	80	0.36	151
219688	98.7	$0.13\pm0.10$	152	0.05	152	0.08	152
220058	1995	$1.17\pm0.20$	67	1.2	65	0.1	121
220116	535	$4.21\pm0.18$	81	2.5	85	1.77	75
223387	1995	$2.11\pm0.20$	67	1.2	65	0.9	70
224544	571	$0.75\pm0.08$	66	0.65	100	0.79	41
224559	417	$0.48\pm0.12$	37	0.55	50	0.24	170
224686	115	$0.08\pm0.20$	86	0.04	88	0.04	84
224905	1510	$1.03\pm0.20$	100	1.5	75	1.15	143
225095	358	$1.27\pm0.20$	79	0.85	65	0.66	93
225985	2100	$0.51\pm0.18$	149	0.5	80	0.9	159
227836	2750	$1.80\pm0.20$	82			1.8$\dagger$	82
228104	1750	$0.50\pm0.20$	70	0.45	160	0.95	70
228438	3500	$1.15\pm0.20$	170	1.0	170	0.15	170
228860	1995	$1.38\pm0.18$	102	0.5	80	1.08	110
229171	1995	$0.78\pm0.20$	20	1.5	180	1.0	75
230579	1200	$1.84\pm0.18$	99	0.7	80	1.35	109
232552	306	$2.21\pm0.20$	117	1.2	110	1.09	125
235795	2180	$0.89\pm0.18$	40	2.0	48	1.2	144
236689	1580	$2.72\pm0.18$	91	2.5	90	0.25	101
236935	1510	$4.29\pm0.18$	91	3	95	1.4	82
237056	780	$5.90\pm0.18$	125	5.5	118	0.9	152
239758	2090	$2.12\pm0.18$	59	1.5	45	1.1	80
244894	2290	$1.61\pm0.18$	19	1	50	1.45	0
245493	1050	$1.11\pm0.18$	165	1	160	0.2	12
246878	2500	$0.55\pm0.18$	133	1	180	1.1	104
248753	395	$2.12\pm0.20$	139	0.9	165	1.7	127
249695	1820	$1.06\pm0.18$	176	1.2	170	0.25	54
250028	1000	$4.29\pm0.18$	164	3.5	175	1.7	138
250163	1450	$1.66\pm0.18$	136	2	165	1.8	100
250289	546	$2.95\pm0.18$	162	1.95	165	1	156
254647	2630	$1.11\pm0.18$	11	0.6	180	0.6	22
256577	1995	$1.61\pm0.18$	148	0.6	170	1.3	140
298298	1900	$1.37\pm0.16$	158	1.5	170	0.6	113
300584	1820	$1.14\pm0.10$	52	0.5	150	1.6	54
302724	1320	$1.40\pm0.10$	59	0.8	80	1	42
302838	2290	$0.97\pm0.10$	157	1.2	150	0.4	40
305382	1900	$1.04\pm0.10$	145	1.5	100	1.8	173
306209	3150	$1.65\pm0.10$	94	1.1	95	0.6	90
306962		$1.80\pm0.10$	80	1.4	100	1.2	55
307350	1820	$0.87\pm0.10$	91	1.2	100	0.5	28
312973	1900	$0.69\pm0.25$	147	0.8	110	0.9	175
316568	1380	$1.65\pm0.08$	13	1.2	180	0.8	35
316587	3000	$0.51\pm0.18$	155	1.2	180	1	125
316589	2550	$1.20\pm0.18$	19	1.2	180	0.8	55
CD-30 5559		$0.92\pm0.10$	171	0.9	150	0.7	25
CPD-60 4551	461	$1.71\pm0.10$	97	1.5	85	0.71	128
CPD-60 4708		$1.42\pm0.10$	73	1.5	80	0.38	25
BD-13 2040		$0.65\pm0.20$	14	0.9	175	0.6	60
BD-13 4936	1250	$1.61\pm0.18$	101	0.9	80	1.1	35
BD+35 1169	1500	$1.66\pm0.18$	133	1.4	150	0.9	105
BD+41 3731	2500	$0.35\pm0.15$	125	0.9	90	0.85	169
BD+47 3302	1650	$1.20\pm0.18$	0	1	20	0.8	155
BD+49 3735	2400	$1.67\pm0.18$	35	1.4	40	0.4	15
BD+53 2964	1650	$1.17\pm0.18$	61	2.0	50	1	125
BD+56 473	1150	$4.38\pm0.18$	111	3.5	120	1.5	90
BD+56 484	2290	$3.46\pm0.18$	124	3.5	120	0.5	170
BD+56 493	2290	$3.64\pm0.18$	126	3.5	120	0.8	165
BD+56 582		$3.69\pm0.20$	106	3.5	120	1.7	70
BD+56 624		$3.13\pm0.20$	107	3.5	120	1.5	61
BD+57 515	1900	$4.29\pm0.18$	111	4	115	0.65	80
BD+57 607		$3.82\pm0.20$	117	4	115	0.3	177
BD+58 458	1500	$4.70\pm0.18$	115	4.5	115	0.2	115
BD+59 2829	3150	$1.38\pm0.18$	74	1.3	80	0.3	40

$\dagger$ - observed polarization.
$\odot$ - no objects in 3$^\circ$-vicinity, $p_{\rm intr}$ from Ghosh et al. (1999).
$\otimes$ - $p_{\rm intr}$ from Ghosh et al. (1999).
$\oslash$ - $p_{\rm intr}$ from McLean & Brown (1978) reduced to V band.

The columns in this table give HD number or other identification, the distance to the object in parsec (marked by italic when it was derived from reddening or from an absolute magnitude-spectral type/luminosity calibration), the value of observed V band polarization with its standard deviation ( $p_{\rm obs}$ in percents and $\theta_{\rm obs}$ in degrees), the estimated value of interstellar polarization ( $p_{\rm is}$ in percents and $\theta_{\rm is}$ in degrees), the calculated value of intrinsic polarization ($p_{\ast }$ in percents and $\theta_{\ast}$ in degrees).