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Astron. Astrophys. 345, 505-513 (1999)
6. UBV photometry
6.1. Transformations
As was already mentioned in Sect. 2, our observations were carried
out from two sites: Mt. Suhora ( ) and
Bial ków ( ). We calculated the
instrumental magnitudes in all bands (separately for each site) with
respect to five relatively bright stars in the field. Next, the
differential magnitudes were corrected for second-order extinction
effects and averaged. These mean values were used in subsequent
transformations.
Because the data carried out at
Suhora were the most numerous, we transformed them first. In the
transformations, we used mainly the photoelectric photometry of
Johnson & Morgan (1955) and Schild (1965). Because of the lack of
red photoelectric standards, we also applied the photographic
photometry of some bright, well-isolated stars from Moffat & Vogt
(1974) giving them half the weight given to photoelectric
measurements. The resulting transformation equations follow:
![[EQUATION]](img70.gif)
![[EQUATION]](img71.gif)
![[EQUATION]](img72.gif)
where numbers in parentheses are the rms errors of the preceding
numbers and the lowercase symbols denote the instrumental magnitudes.
The internal standard deviations are 0.031, 0.044, and 0.059 mag, for
Eqs. (6), (7), and (8), respectively. The above three equations were
used to transform only the Suhora measurements. Then, using
magnitudes and
colour indices obtained in this way,
we transformed Bial ków
photometry with equations:
![[EQUATION]](img75.gif)
![[EQUATION]](img76.gif)
This time, the lowercase letters stand for Bial ków
instrumental magnitudes. The standard deviations are equal to 0.024
and 0.017 mag for Eqs. (9) and (10), respectively. Subsequently,
for stars observed in both sites, transformed V magnitudes and
colour indices were combined
together. Resulting standard
magnitudes and colour indices for
184 stars as well as colour indices
for 79 stars are given in Table 4. The full version of this table
is available only in electronic form at CDS via anonymous ftp to
130.79.128.5. Here - for reader's
convenience - we present only the data for 10 variable stars.
![[TABLE]](img90.gif)
Table 4. photometry of stars in h Persei. The columns are: (1),(2), X and Y coordinate in Fig. 1, (3), Oosterhoff (1937) number, (4), Wildey (1964) number, (5) Moffat & Vogt (1974) number, (6), magnitude, (7), colour index, (8), reference to photometry (`S' means that the star was observed from Suhora only, `B', from Bial ków only, `S+B', both from Suhora and Bial ków. The photometry transformed with the use of I-filter measurements [Eqs. (11)-(14)] are additionally flagged with `I'), (9), colour index, (10), right ascension (epoch 2000.0), (11), declination (epoch 2000.0), and (12), remarks.
The -filter observations were not
transformed to the standard Cousins
system because there are no
standards in the field. However, because the I-filter
photometry is the deepest, we incorporated these measurements in
another way, namely using them for the transformations for faint stars
for which the photometry was not
reliable. The transformation equations
![[EQUATION]](img92.gif)
![[EQUATION]](img93.gif)
for Bial ków and
![[EQUATION]](img94.gif)
![[EQUATION]](img95.gif)
for Suhora give us additional
photometry for 74 faint stars. The standard deviations for
Eqs. (11)-(14) are, in succession, equal to 0.014, 0.024, 0.004, and
0.025 mag. This photometry is also included in Table 4. One
may doubt the reality of the colour
indices transformed from indices if
red and reddened stars are used in transformation, but in our case
both the stars used to obtain the transformations and those
transformed later are mostly reddened cluster members. It follows that
the systematic effects can be important only for a few faint
non-members.
Because the mean instrumental magnitudes were calculated from all
search frames, the internal accuracy of our
photometry is very good - it is
better than 1 mmag for brightest stars and reaches about
0.02 mag for the faintest. Obviously, owing to the transformation
errors, the magnitudes given in Table 4 are much less
accurate.
Magnitudes and colours of eclipsing variables in Table 4 were
calculated for the phases of maximum light. For remaining variables
mean magnitudes and colours calculated from all our observations are
given. Equatorial coordinates given in the tenth and eleventh column
of Table 4 were derived using the Guide Star Catalogue (GSC)
positions of 57 stars in the observed field. Regarding the accuracy of
the stars' positions in GSC, the coordinates given in Table 4 are
accurate to within 0 1.
6.2. Colour-magnitude and colour-colour diagrams
The cluster colour-magnitude (CM) and colour-colour diagrams are
shown in Fig. 9. photometry of faint
stars obtained with the use of
indices is shown with plus signs in the CM diagram.
![[FIGURE]](img113.gif) |
Fig. 9. a The colour-magnitude diagram for the observed field in h Persei. The two Cephei stars are denoted by open, and the other stars by filled circles. Variables are enclosed in large open squares. Stars for which photometry was transformed from the instrumental measurements are shown with plus signs. b Colour-colour diagram for stars which photometry is available. Our photometry is shown with the same symbols as in the left diagram. Crosses are points in which the indices were taken from other sources, photoelectric (large symbols) or photographic (small symbols). The intrinsic colour-colour relation for dwarfs, shown with dashed line, was taken from Caldwell et al. (1993). The same relation for the mean value of reddening, = 0.52 mag (long solid line), and = 0.47 and 0.57 mag (short solid lines) are also plotted.
|
As in Persei, both the h
Persei Cephei stars, Oo 692 and
Oo 992, lie close to the cluster turn-off point (Fig. 9a). The
positions of these two stars bracket a few other stars, including the
Eri star Oo 922. The V
magnitudes of the Cephei stars (9.35
for Oo 692, and 9.90 for Oo 992) are in the same range as in
Persei. The scatter in the cluster
main sequence is real and is a result of small differences in
reddening. The leftmost star in the CM diagram is Oo 622. The
colour index for this star is equal
to -0.53 (Tapia et al. 1984). This means that it is probably an
early-type star, either slightly less reddened than the other cluster
stars or simply a foreground object.
In order to derive the average reddening of the cluster, we moved
the intrinsic
vs. relation for dwarfs taken
from Caldwell et al. (1993) (dashed line in Fig. 9b) along the
reddening line with the average parameters derived by Turner (1989),
i.e.,
![[EQUATION]](img116.gif)
The best agreement was obtained for
= 0.52 mag (long solid line in
Fig. 9b). As was shown by Turner (1989), the slope of the reddening
line is not unique throughout the sky. This, however, does not greatly
affect our result: even if the extreme values of the slope of
reddening line derived by Turner (1989) are assumed, the best-fit
value of for h Persei differs
from 0.52 mag by no more than 0.01 mag. The average value of
is in general agreement with previous
determinations (see Tapia et al. 1984and references therein; Pandey et
al. 1989; Natali et al. 1994). As can be seen in Fig. 9b, most of the
cluster stars have reddenings which differ from the mean value by no
more than 0.05 mag. This range of individual reddenings is smaller
than obtained by Wildey (1964), but this disagreement will become
understandable in the view of the accuracy of his photometry (see next
subsection).
Unfortunately, no Strömgren photometry is available for the
two Cephei stars we discovered.
Therefore, we could not convert Strömgren indices to
and compare their positions with
those in other open clusters, as we did in Paper I (see Fig. 9 there
in). It seems, however, rather certain that both these stars are still
in the core-hydrogen burning phase.
6.3. Comparison with previous work
We compared our photometry with
previous photoelectric (Johnson & Morgan 1955; Schild 1965; Tapia
et al. 1984) and photographic (Wildey 1964; Moffat & Vogt 1974)
studies. The result of these
comparisons is shown in Fig. 10, and the mean differences are given in
Table 5.
![[TABLE]](img130.gif)
Table 5. Mean differences between our and other photometric studies. SD stands for the standard deviation from the mean difference which is given in the preceding column. N is the number of stars in common with a given author. Both the mean differences and SD are expressed in mag.
Notes:
a The mean differences, and , and the corresponding standard deviations, were in each case calculated without the one star that deviated most. This was Oo 837 in the and Oo 885 in the .
Our photometry agrees quite well with all photoelectric studies.
Out of the two photographic data
sets, the one of Moffat & Vogt (1974) is much better. The
differences between our and Wildey's (1964) V measurements are
mostly positive and have very large scatter, even for bright stars. On
the other hand, his indices are
systematically larger than ours, and have a large scatter too. This
means that this photometry is very uncertain, as was already claimed
by Tapia et al. (1984). Consequently, Wildey's (1964) conclusions
concerning the cluster parameters and reddening should be viewed with
caution. The photographic photometry
of Moffat & Vogt (1974) have a much smaller scatter than Wildey's,
and agrees better with the photoelectric measurements. This is why
some of the stars measured by them were used in our
transformations.
© European Southern Observatory (ESO) 1999
Online publication: April 19, 1999
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