A&A 376, 144-153 (2001)
DOI: 10.1051/0004-6361:20010974
A. Pigulski - G. Kopacki - Z. Koaczkowski
Wrocaw University Observatory,
Kopernika 11, 51-622 Wroc
aw, Poland
Received 13 December 2000 / Accepted 14 June 2001
Abstract
We present new
H
photometry of the
young open cluster NGC663. The H
photometry is complete
down to magnitude
,
corresponding to spectral type
A5 for the cluster members. This allows detection of mild and strong
H
emission in all B-type stars in the cluster. In addition to
the 22 Be stars known in the observed field of NGC663, we discovered
four new faint stars of this type. We find that Be stars in NGC663
cover the whole range of the B spectral type. They are, however, most
populous among stars with spectral types falling in the range between B0 and B3, where their fraction amounts to 31
8%. Among B-type
stars later than B3, Be stars are much less abundant: only 7 out of
101 observed stars, that is, 7
3%, were detected. About 70%
of the observed Be stars in NGC663 show detectable variations of
light. In the time interval covered by our observations, the ranges
of the largest variations reach 0.4 mag in the
band. By
means of the isochrone fitting, we derived the cluster distance of 2.1 kpc, age of 20-25 Myr, and the mean colour excess
=
0.54 mag, with a
0.1 mag scatter due to differential reddening.
Key words: stars: emission-line, Be - open clusters and associations: individual: NGC663
The differences in Be-star fractions in open clusters are still not
well understood. Some investigators postulate age (Fabregat &
Torrejón 2000) and metallicity (Maeder et al. 1999) effects. In order to study this problem, the Be
star fractions have to be known in many clusters and in the whole
range of B spectral types. Unfortunately, the observed fractions of
Be stars are very often systematically biased because the searches are
not complete. This is mainly true in the case of spectroscopic
searches which, for practical reasons, were usually confined to the
brightest and most well-isolated stars in a cluster. As a
consequence, Be star fraction for late B-type stars was often
underestimated. The situation was improved when the era of CCD
H
photometry began. Although with this technique faint
emission cannot be detected, it works well in crowded fields and needs
much less telescope time than does spectroscopy. Even with a
small-size telescope, late B-type stars in many open clusters can be
easily reached, so that complete information on the Be star content
can be obtained.
Detection of Be stars by means of the CCD photometry is usually
carried out with a filter centred on H
and two other filters
located in the Paschen continuum (see, e.g., Grebel 1997).
It is even better when two filters of different width centred on
H
are used, because they may define a photometric index
independent of the interstellar extinction (see, e.g., Goderya &
Schmidt 1994). Such a pair of H
filters was used in
our observations.
NGC663 is a young galactic open cluster very rich in Be stars. According to Sanduleak (1979), it contains over twenty Be stars, that is, more than one third of all its early B-type stars. This is why this cluster was chosen as one of the first targets in our Be star survey among northern open clusters. Within this search, we have recently discovered that NGC7419, another open cluster of similar age, contains over thirty Be stars (Pigulski & Kopacki 2000). Our project is also complementary to the ongoing search for variables of type B in young open clusters (see Pigulski et al. 2000 and references therein).
All observations of NGC663 were carried out at the Biaków
Observatory of the Wroc
aw University with a 60-cm telescope
equipped with a
pixel CCD camera and an autoguider.
Since the camera field of view is a 6
4
rectangle, 15 overlapping fields were covered in order to observe most
of the cluster stars. The borders of the observed field, called
hereafter the "P field'', are shown with a solid line in
Fig. 1, while a detailed map, showing all stars we measured,
is given in Fig. 2.
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Figure 1:
A 30![]() ![]() ![]() ![]() ![]() |
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Figure 2:
Schematic view of our fifteen 6![]() ![]() ![]() ![]() |
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The observations in each field consisted of a series of frames taken
through filters reproducing the
bands of the
Johnson-Cousins (Johnson & Morgan 1953; Cousins
1976) system and a single frame taken through the narrow
H
filter. The latter was usually sandwiched between two frames
taken through the wide H
filter. The observations were carried
out during 10 nights between March 7, 1997 and September 1, 1998.
Typical exposure times were the following: 2000 s for narrow
H
,
600 s for wide H
and B, and less than 200 s for
the remaining filters.
NGC663 was also observed in 1999 in the above-mentioned program of
searching for variable B-type stars in the northern open clusters.
During these observations, however, only four slightly overlapping
fields were observed in V and
bands. The area covered
during this search is shown in Fig. 1 with dashed lines and
will be called hereafter the "V field''. The results of the variability
survey in the V field will be published separately (Pigulski et al. 2001). In this paper we present only that part of these
results which refer to the variability of Be stars (see Sect. 6).
The CCD frames were calibrated in a standard way (see, e.g., Jerzykiewicz et al. 1997) and then reduced using the point-spread-function fitting in the Daophot II package of Stetson (1987).
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(1) |
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(2) |
Since till now no
photometry was published for
NGC663, in order to transform our observations in these bands to the
standard system, we observed NGC663 on three photometric nights in
September and October 2000 together with another open cluster,
NGC7790. We have chosen NGC7790 instead of the usually observed
Landolt (1992) fields for two reasons: (i) this cluster is
much closer to NGC663 than the Landolt (1992) standards,
(ii)
standard photometry of NGC7790 has been recently
published by Stetson (2000).
Using the photometry of NGC7790 provided by Stetson (2000),
we derived the following transformation equations:
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(3) |
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(4) |
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Figure 3:
Comparison of BV photometries. All differences, ![]() ![]() |
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The photometry we provide in Table 2 is not homogeneous in the sense that the central regions were observed more frequently than those situated at the borders of the P field. In consequence, stars of the same magnitude, but from different regions, may have photometry which differs in accuracy by a factor of up to 5. For this reason, the photometric errors are also given in Table 2.
Having transformed our BV photometry to the standard system, we compared it with the following previous photometric studies: CCD measurements of Phelps & Janes (1994), photoelectric photometry of Hoag et al. (1961), photographic data of the same authors as well as those of McCuskey & Houk (1964), Moffat & Vogt (1974), and van den Bergh & de Roux (1978). Results of this comparison are shown in Fig. 3.
Photoelectric measurements of Hoag et al. (1961) were the primary source of standard stars used in our transformations. Hence the scatter in panel A of Fig. 3 is small. Out of four photographic photometries, the best agreement and the smallest scatter are shown by the photometry of van den Bergh & de Roux (1978). The differences between our photometry and the photographic photometries of Moffat & Vogt (1974) and McCuskey & Houk (1964) have larger scatter, while that of Hoag et al. (1961), especially the B photometry, exhibits a clear systematic, magnitude-dependent effect.
Surprisingly, the CCD photometry of Phelps & Janes (1994)
also differs systematically from our measurements. While for stars
with magnitudes in the range between 12 and 16 in B and V the
agreement is quite good (apart from a 0.06 mag shift in B), for
stars fainter than 16 mag, the photometry of Phelps & Janes
(1994) is systematically fainter and shows a large scatter.
In addition, stars brighter than 12 mag are also systematically
fainter in their photometry. The same effect can be seen in both Band V, reaching 0.5 mag for stars of the 10th magnitude.
The field of NGC663 is one of the most populous as far as Be stars
are concerned. In the following discussion of Be stars in this open
cluster, we shall confine ourselves to the 30
30
area shown in Fig. 1.
The first Be star discovered in this area was MWC428 = BD+60325 (Merrill & Burwell 1943). A few years later, the next
four Be stars: MWC698 = BD+60
332A, MWC700 = BD+60
340, AS42 = BD+60
341, and AS43 = BD+60
343B
were found by the same authors (Merrill & Burwell 1949,
1950). Since subsequent observations did not detect emission
in MWC700 (known also as Sanduleak 11), its Be nature is not certain.
As was pointed out by Sanduleak (1979), it could happen that
the BD number was incorrectly assigned by Merrill & Burwell
(1949) and the star is identical with AS42.
González & González (1954) were the first to show that the cluster contains a large number of Be stars. They found 13 new stars of this type in the area under investigation. Subsequent studies (Dolidze 1975; Schild & Romanishin 1976; Coyne et al. 1978; MacConnell & Coyne 1983; Sanduleak 1990) added seven more stars, increasing the number of certain Be stars in NGC663 to 24. These stars are listed in the compilation of Kohoutek & Wehmeyer (1997).
Sanduleak (1979) has provided a list of 27 Be stars in NGC663, extending it by two objects in a subsequent paper (Sanduleak 1990). Since the numbers from Sanduleak's lists are widely used in the literature, we shall also refer to these stars preceding the number with "Sanduleak''. We now know, however, that not all Sanduleak's stars are still considered to be Be stars. Sanduleak 7, 18, and 19, announced to be Be stars by Coyne et al. (1978), appeared to be non-emission stars after revision (MacConnell et al. 1983). In addition, two other stars, Sanduleak 15 and 24, were found to have doubtful or very weak emission (Schild & Romanishin 1976; Voight 1965, respectively). Later observations (Sanduleak 1979, 1990; Torrejón et al. 1997; this paper) did not reveal emission in these stars either. Thus, only 23 stars with Sanduleak's numbers should be considered as certain Be stars and one, Sanduleak 11, as a probable Be star. We list these stars in Table 1. This table includes also five other emission-line stars: four new Be stars we discovered in the P field and star GG90 = D01+030 (see caption of Table 1 for designations) which probably is not a member, but is situated in the field shown in Fig. 1.
S | W | HBH | X | Y | Right Asc. | Decl. | ![]() |
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Other designations |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10) |
-- | 240 | -- | 1056.6 | 309.1 | 1 45 11.65 | +61 10 22.4 | 12.972 | 2.034 | new Be star |
1 | 243 | 6210-32 | 990.8 | 638.9 | 1 45 18.02 | +61 06 55.9 | 9.902 | 1.463 | BD+60![]() |
-- | 417 | -- | -- | -- | 1 45 18.31 | +60 58 09.2 | -- | -- | GG90, D01+030 |
2 | 210 | 6210-54 | 759.9 | 544.1 | 1 45 37.77 | +61 07 59.2 | 12.054 | 1.330 | GG93, VES610, D01+031 |
3 | 130 | -- | 727.3 | 65.1 | 1 45 39.63 | +61 12 59.6 | 12.094 | 1.388 | GG94, VES611 |
4 | 170 | -- | 657.6 | 415.8 | 1 45 46.35 | +61 09 20.9 | 11.771 | 1.977 | GG95 |
-- | 128 | -- | 622.8 | -259.1 | 1 45 47.85 | +61 16 23.6 | 13.122 | 1.949 | new Be star |
5 | 93 | 6210-56 | 537.6 | 91.8 | 1 45 56.07 | +61 12 45.3 | 11.585 | 1.444 | GG97, VES613 |
6 | 92 | 6210-57 | 500.6 | 92.2 | 1 45 59.27 | +61 12 45.6 | 10.230 | 1.982 | BD+60![]() |
28 | 84 | 6210-34 | 463.4 | -125.1 | 1 46 02.03 | +61 15 02.1 | 13.535 | 1.333 | |
8 | 53 | 6210-15 | 419.8 | 8.5 | 1 46 06.09 | +61 13 39.1 | 11.115 | 1.388 | BD+60![]() |
29 | 51 | 6210-73 | 328.6 | 3.3 | 1 46 13.99 | +61 13 43.8 | 12.846 | 1.820 | |
9 | 10 | 6210-31 | 255.8 | -55.6 | 1 46 20.19 | +61 14 21.6 | 11.442 | 1.741 | GG98, VES617 |
-- | 61 | -- | 215.7 | 304.1 | 1 46 24.37 | +61 10 37.3 | 13.159 | 1.610 | new Be star |
13 | 141 | 6210-13 | 194.0 | 585.1 | 1 46 26.80 | +61 07 41.8 | 10.197 | 1.341 | BD+60![]() |
10 | 8 | -- | 178.8 | -37.1 | 1 46 26.87 | +61 14 11.2 | 11.905 | 1.999 | GG99 |
11 | 107 | -- | 186.3 | 361.3 | 1 46 27.02 | +61 10 02.0 | 9.892 | 2.130 | BD+60![]() |
12 | 21 | -- | 174.1 | 132.2 | 1 46 27.64 | +61 12 25.5 | 10.341 | 1.725 | |
14 | 6 | 6210-36 | 135.1 | -64.9 | 1 46 30.60 | +61 14 29.3 | 11.656 | 1.446 | VES620 |
16 | 39 | -- | 75.7 | -189.1 | 1 46 35.49 | +61 15 48.0 | 9.851 | 1.680 | BD+60![]() |
17 | 2 | 6210-14 | 79.9 | 16.6 | 1 46 35.57 | +61 13 39.2 | 11.857 | 1.374 | GG101 |
20 | 288 | 6210-59 | -- | -- | 1 46 57.53 | +61 01 40.4 | -- | -- | GG103 |
21 | 121 | 6210-51 | -194.5 | 134.2 | 1 46 59.55 | +61 12 29.8 | 13.012 | 1.565 | VES624 |
22 | 194 | 6210-50 | -218.3 | -720.0 | 1 47 00.17 | +61 21 23.7 | 11.039 | 1.837 | GG104, VES625 |
23 | 124 | 6210-33 | -252.3 | -349.0 | 1 47 03.70 | +61 17 32.2 | 12.072 | 2.050 | D01-034 |
-- | 151 | -- | -402.9 | 61.2 | 1 47 17.46 | +61 13 18.2 | 13.870 | 1.859 | new Be star |
25 | 222 | 6210-35 | -501.4 | 499.3 | 1 47 26.72 | +61 08 44.8 | 11.264 | 1.963 | GG108, VES628 |
26 | 224 | 6210-25 | -663.8 | -419.7 | 1 47 39.34 | +61 18 20.7 | 11.908 | 1.426 | GG109, VES630, D01+036 |
27 | 297 | 6210-26 | -- | -- | 1 48 23.04 | +61 15 53.2 | -- | -- | GG110 |
The H
photometry we made was used to find stars showing
H
emission. This was done by means of an
index,
defined as a difference between the magnitudes of a star through a
narrow and a wide filter. Our filters and, consequently, the
index resembles those used by others, namely Abt & Golson
(1966), Tebbe (1969), Feinstein (1974),
Dachs & Schmidt-Kaler (1975), Claría & Escosteguy
(1981), Strauss & Ducati (1981), and Goderya &
Schmidt (1994). We made, however, no attempt to transform
our
to that defined by others; it was left in the
instrumental system. In Fig. 4 we show the
index
for 442 stars with
15.4 mag, the limiting magnitude of
our H
photometry. The same stars are shown in
Fig. 5, where the
index is plotted as a function of
.
Out of the 25 known Be stars in the area shown in Fig. 1, 22 fall within the P field, and 14 within the V field. As can be seen in
Fig. 4, except for doubtful Be star Sanduleak 11, all these 22 known Be stars indeed show emission, that is, have
placing
them above the locus of cluster non-emission stars, shown with a
continuous line. The line was obtained from the spectra given by
Danks & Dennefeld (1994) which were combined with the
transmission curves of the H
filters provided by the
manufacturer. Except for known Be stars, some other stars are found
above this line as well. Because some foreground late-type stars can
also have weak emission, we need to separate them from the Be stars
first. For this reason, we made a rough selection of non-members with
the use of the cluster CM diagram (see Sect. 7). Stars selected as
non-members are plotted with crosses in Figs. 4 and
5. Obviously, with this kind of selection, we cannot
indicate all non-members which contaminate the cluster main sequence.
Many objects with
14 mag in Fig. 4 situated
clearly above the continuous line are certainly such unrecognized
non-members.
Using Fig. 4 we can conclude that four stars, namely W61,
128, 151, and 240 (we precede star numbers taken from Wallenquist
(1929) with "W''), may be classified as new Be stars. A
star was selected as a new Be star if: (i) the difference between
the observed
and the value taken for a star of the same
magnitude showing no emission (continuous line in Fig. 4)
was negative and its absolute value was larger than 4
,
where
is the rms error of
,
and (ii) the star
was not selected as a non-member.
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Figure 4:
The ![]() ![]() ![]() ![]() ![]() |
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Figure 5:
The ![]() ![]() |
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All four new Be stars are among the faintest Be stars in NGC663,
implying that the previous Be star surveys were rather complete for
the brightest cluster stars. It can be also seen in Fig. 4
that Be stars with strong and intermediate H
emission occur in
the whole range of the B spectral type. Most of them, however, occupy
the range between B0 and B3. Maeder et al. (1999) compared
the fractions of Be stars in clusters of an age comparable to that of
NGC663. In the widest interval considered by these authors,
-5 <
MV < -1.4 mag, they found that NGC663 contained 34
11% of Be stars. The above-given interval of MV corresponds
to
mag. There are 59 stars we observed in
this interval of
,
including 18 Be stars. This gives the
revised fraction of Be stars in NGC663 equal to
%, with
the error calculated assuming Poisson statistics.
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Figure 6:
The ![]() |
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Among stars fainter than
mag, we find 7 Be stars,
including all four discovered in our search. Comparing this number
with the number of all B-type stars later than B3 (101), we obtain the
fraction of
%.
For 13 Be stars, simultaneous (in the sense that a narrow Hframe was sandwiched between two wide H
ones) H
photometry was made in one epoch only. For the remaining 13 Be stars,
such observations were made on two different epochs allowing detection
of possible changes in the emission strength. The largest change of
we detected was 0.045 mag for Sanduleak 3 and two epochs
separated by about 550 days. The epochs of H
observations and
corresponding
values are given in Table 3, available in
electronic form only.
The
indices given in Table 1 were calculated using all frames,
including some additional wide H
frames made on epochs other
than those given in Table 3. For this reason the
indices
given in Table 1 (average values) are not exactly the same as those
presented in Table 3 (epoch photometry).
As we mentioned above, 14 Be stars from Table 1 are situated in the V field, the target of the variability survey. Out of these 14 stars, 10 were found to be variable (Pigulski et al. 2001) and only four
(Sanduleak 11, 17, 28, and the newly discovered Be star, W128) show
no light variation above the detection limit. Nine variable Be stars
show aperiodic or quasiperiodic variatons with ranges up to 0.4 mag in .
Only Sanduleak 10 is a periodic variable of
Eridani type with a period of 0.67298 d (Pigulski et al. 2001). Two drops of brightness seen in the light curve of
Sanduleak 29 on HJD2451200 and 2451439 could be eclipses. A
large percentage of variables among Be stars is not unusual, because
virtually all these stars exhibit some degree of photometric
variability on different time-scales. The light curves of all variable
Be stars situated in the V field are shown in Fig. 6.
Although observed on only two or three nights, some of the Be stars
from outside the V field can also be classified as variables. These
are: Sanduleak 1 (
mag), 2 (
0.09 mag), 3 (
0.10 mag),
4 (
0.10 mag), and 23 (
0.05 mag). The remaining Be stars, that is, Sanduleak 11, 21,
22, 25, 26, W128, 151, and 240, show no clear evidence of variability
exceeding the photometric errors of our photometry.
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Figure 7:
CM diagram for NGC663. Crosses denote stars identified as
non-members (see Sect. 7). Diamonds are used to indicate the
emission-line objects. The dashed line shows the limit of the
H![]() |
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In order to plot an average position of the variable Be star in the CM diagram (Fig. 7), one should use observations made
simultaneously in all bands. Such observations were chosen for
averaging, but the epochs of the averaged magnitudes are not the same
for all Be variables (different fields were observed on different
nights). These average epochs of the
photometry are
the following: HJD2450515.3 for Sanduleak 9, 14, and 16;
HJD2450518.4 for Sanduleak 13; HJD2450813.3 for Sanduleak 1, 2, and 4; HJD2450845.3 for Sanduleak 23; HJD2450846.3
for Sanduleak 5, 6, 8, 12, and 29; and HJD2451058.6 for
Sanduleak 3. Observations of non-variable Be stars and Sanduleak 10
(
Eridani-type variable) were averaged using all available
photometry.
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Figure 8:
Dereddened
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The members of NGC663 form a well-defined main sequence in the CM diagram (Fig. 7). The main sequence is slightly widened by a differential reddening. Some non-members can be seen redward of the cluster main sequence. Non-members also contaminate the lower part of the main sequence, because with the mean reddening of NGC663 (0.8 mag in terms of the E(B-V) colour excess according to Phelps & Janes 1994), foreground late-type dwarfs and the reddened cluster B- and A-type stars coincide in the CM diagram.
Since the late-type non-members also affect Fig. 4, we
would like to distinguish them from the cluster members. Although not
unambiguous, the only selection we can practically perform is the
selection that uses two-colour photometry. For instance, the
non-members can be roughly selected by a dereddening procedure which
takes into account the information on the spatial distribution of
stars. Assuming that there is no reddening within the cluster and
adopting a certain value of the spatial scale of reddening variation,
characterized by a parameter
(see below), a reddening
map can be derived and then used to correct the colours for
differential reddening and magnitudes for differential extinction.
We made an attempt to calculate the variable part of the reddening
with the procedure applied by Pigulski & Koaczkowski
(1998) to the Cygnus OB2 association. We refer the reader to
that paper for details of the calculations. Unlike that case,
however, for NGC663 we did not assume a priori the shape of the
unreddened main sequence. Instead, the sequence was fitted by a third
degree polynomial. The procedure was done iteratively in the
following steps: (i) fitting the cluster main-sequence with a
polynomial, (ii) calculating the differences between the observed
position of a star and the intersection of the reddening line passing
through that point with the main sequence (the difference was measured
along the reddening line), (iii) calculating the reddening map with
the use of the above differences transformed into reddenings and then
deriving the residuals from this map for each star, (iv) selecting of
non-members. In the last step, a star was identified as a non-member
if its residual from the reddening map was larger than an assumed
threshold. The rejected stars were not used in fitting the cluster
main sequence in the next iteration. The iterations were terminated
when no change in residuals occured.
The appearance of the reddening map depends on the assumed value of
(see Pigulski & Ko
aczkowski 1998 for
the definition of this parameter) which, generally, is not known. The
lower value of
results in a more detailed extinction map
and narrower dereddened main sequence, a larger value gives the
opposite effect. Fortunately, the selection of non-members is almost
independent of the choice of
.
The main parameter which
determines how many stars will be selected as non-members is the
threshold adopted in step (iv) of the dereddening procedure. We have
assumed the threshold to be equal to 0.1 mag. Despite some
subjectivity in the method of selecting non-members, we think it is
reasonable to use it, because it helps to find new Be stars. In
Fig. 4 we see that some late-type non-members and Be stars
with weak emission have similar values of
.
The selection we
made allows us to distinguish fairly well between these two groups of
stars.
Using the extinction map derived in Sect. 7, we corrected the CM diagram for the differential reddening. There is an advantage in using such a CM diagram for the derivation of cluster parameters, especially its age. The problem has already been pointed out by Phelps & Janes (1994): a large spread in colours of stars close to the turn-off point led them to the conclusion that the possible range of age of NGC663 is quite wide, between 12 and 25 Myr. An age of about 10 Myr was also proposed by Liu et al. (1991) and Tapia et al. (1991), while Mermilliod & Maeder (1983) include NGC663 in the group of clusters with an age of about 21 Myr.
We used our CM diagram corrected for differential reddening for the
estimation of age, true distance modulus and the mean reddening of
NGC663. For this purpose we applied the isochrones of Bertelli et al. (1994) for Y = 0.28 and Z = 0.02. The results of
fitting the isochrones are shown in Fig. 8. We finally
adopt 11.6 mag as the true distance modulus of NGC663 (it
corresponds to a distance of 2.1 kpc). This agrees very well with the
recent determinations of the distance modulus given by others:
11.80 mag (Moffat 1972), 11.55 mag (van den Bergh & de Roux
1978), 12.03
0.20 mag (Tapia et al. 1991),
12.25 mag (Phelps & Janes 1994), and 11.6
0.7 (Fabregat
& Torrejón 2000). The mean cluster colour-excess is
= 0.54 mag. This corresponds to
0.83 mag, in good agreement with other determinations (Johnson et al. 1961; Hoag 1965; Moffat 1972; Phelps &
Janes 1994). In our fit we assumed the ratio of total to
selective absorption
= 3.7.
As can be seen in Fig. 8, the cluster upper main sequence
is best fitted by the
(time in years) = 7.3 and 7.4 isochrones.
This results in the cluster age in the range between 20 and 25 Myr.
The results of our H
observations indicate that previous
searches detected virtually all bright Be stars in NGC663. As could
be expected, our search resulted in new detections of H
emission in late B-type stars only. The search was complete down to
the A5-type stars in the cluster and revealed that Be stars are most
populous in the range corresponding roughly to the spectral types
between B0 and B3 (see Fig. 4). Although Be stars were also
found among cluster members with the spectral type later than B3,
their fraction amounts to only 7
3%. Such a distribution of Be
stars with spectral type was already known (Schild & Romanishin
1976), but was derived from a small sample of stars and
searches which were not complete. The case of NGC663 confirms the
dependence of Be star distribution as a function of spectral type, but
more clusters of different ages need to be observed in order to come
to a reliable conclusion as to this effect. H
observations
of other northern clusters rich in Be stars are now underway.
In terms of the
colour-index, Be stars in NGC663
are, on average, 0.04 mag redder than non-emission stars which define
the cluster main sequence. This additional reddening, presumably of
circumstellar origin, is observed in almost all open clusters
containing Be stars.
Acknowledgements
We are indebted to Prof. M.Jerzykiewicz for critical reading the manuscript. We also thank the referee, Dr. J.Zorec, for his valuable comments. This research was supported by the KBN grant No. 2P03D2909.