B-type pulsators in the central region of NGC 869 (h Persei)

Transcription

1 Astron. Astrophys. 345, (1999) ASTRONOMY AND ASTROPHYSICS -type pulsators in the central region of NGC 869 (h Persei) J. Krzesiński 1,2, A. Pigulski 3, and Z. Koĺaczkowski 3 1 Mt. Suhora Observatory, Cracow Pedagogical University, Podchora żych 2, PL-3-84 Kraków, Poland (sfkrzesi@cyf-kr.edu.pl) 2 Department of Astronomy, University of Cape Town, Rondebosch 77, Cape Town, South Africa 3 Wrocĺaw University Observatory, Kopernika 11, PL Wrocĺaw, Poland (pigulski; kolaczk@astro.uni.wroc.pl) Received 21 December 1998 / Accepted 4 February 1999 Abstract. In the course of the search for -type pulsators in the central region of h Persei, we discovered two β Cephei stars, Oo and Oo 992, and one SP star, Oo 893. The first two stars are monoperiodic pulsators with periods of and d respectively, and semi-amplitudes smaller than.1 mag. The light curve of Oo 893 can be described by a single day period, slightly non-sinusoidal in shape, and having semi-amplitudes from 19 mmag in I to 48 mmag in U. Oo 893 is the first SP star found in h and χ Persei. In addition, we discovered seven other variables, including three eclipsing binaries and one λ Eri star. One of the binaries is a W UMa-type star and a likely cluster member. We also present new U photometry for 258 stars in the field. The average reddening, estimated from the cluster colourcolour diagram, amounts to E( ) =.52 mag. A.1 mag dispersion of reddenings within the cluster is also seen. Key words: stars: binaries: eclipsing stars: oscillations stars: variables: general Galaxy: open clusters and associations: individual: NGC 869 =hper 1. Introduction This is the second paper presenting the results of the variability search among early -type stars in the central regions of NGC 869 and NGC 884, the young, twin cluster in Perseus, better known as h and χ Persei. In the first paper (Krzesiński & Pigulski 1997; hereafter Paper I), the results of the variability search in the central region of χ Persei were presented. In that paper, we announced the discovery of two β Cephei stars, Oo 2246 and Oo 2299, as well as the nine other variables. These discoveries convinced us that the CCD technique we applied is indeed suitable for finding low-amplitude variables in both clusters. In this paper, we show the results of a similar search in the central region of the other cluster, h Persei (NGC 869). As in χ Persei, the observations were carried out at Mt. Suhora Observatory from 1994 to In 1997, additional observations were Send offprint requests to: A. Pigulski 1 The Oo numbers are following Oosterhoff (1937). made at the iaĺ ków Observatory. The obervations and reductions are briefly described in Sect. 2; in the following three sections we present -type pulsators, eclipsing binaries and other variables discovered during the search. oth the U data and the H-R diagram of the cluster are presented in Sect. 6. In the last section the consequences of our findings are discussed. 2. Observations and reductions As mentioned above, the CCD observations of h Persei started at the Mt. Suhora Observatory in 1994 and were continued until 1997 using mainly Johnson and bands. In addition, observations in the U and I bands were obtained from time to time. In 1997, I observations of this cluster were also collected at the iaĺ ków Observatory. At both sites, 6-cm Cassegraintype telescopes equipped with Photometrics Star I CCD cameras were used. This equipment gives us a 4 6 field of view. The observed field covered the central region of h Persei. Owing to the different orientations of the CCD cameras at both sites and the fact that the observed field was changed slightly several times during the run, the total field covered by our observations was larger than that of a single field. This can be seen in Fig. 1, where all 311 stars we detected are shown. During four seasons between October 18, 1994 and January 1, 1998 we acquired 24, 57, 13, and 1 hours of observations in the U,,, and I bands, respectively. All these data were corrected for bias, dark and flat-fields in the usual way and then reduced with the Daophot ii package (Stetson 1987). In order to improve the signal-to-noise, some consecutive frames in the and I bands were summed prior to reduction. Next, because of variable quality from night to night, some data were binned, so that each point on the light curve has approximately the same weight. UI light curves were obtained for all stars with reasonable photometry with respect to the nearby comparison stars. At this stage of reduction, the data were corrected for second-order extinction effects. Then, they were analyzed for variability using error and periodogram analysis in addition to visual inspection of the light curves. Out of 311 stars detected in the field, 1 were found to be variable. They are listed in Table 1. The light curves of all the variables are available in electronic form from CDS in Strasbourg via anonymous ftp to

2 56 J. Krzesiński et al.: -type pulsators in the central region of NGC 869 (h Persei) W Y X 1 Fig. 1. A schematic view of the observed field in h Persei. Except for W 49, the variables are labeled with their Oosterhoff (1937) numbers. Note that only these stars which we detected are plotted. North is up, east to the left. Table 1. ariable stars in the central region of h Persei. W 49 stands for the star number 49 of Wildey (1964). Range Oo Period P in Type of number [days] [mag] variability β Cephei SP β Cephei λ Eri :.15 Unknown EA (NS 779) Irregular ( 52 Per) or P /2.15 Ellipsoidal? :.53 EA W W UMa 3. -type pulsators From our point of view, we regard the discovery of three pulsating -type stars as of the greatest interest. As it is shown in Table 1, there are two β Cephei stars and one slowly pulsating -type (SP) star in the observed field. We shall describe them now in detail Oo 692 = D The MK spectral type of Oo 692 ( ) was derived by C. oehm and R. Stalio as cited by Franco et al. (1985). Shortly after the beginning of our observations, the star was identified as a possible β Cephei variable, pulsating with a period of about.172 d (Krzesiński 1995). The analysis of the whole data set confirmed the presence of the.172 d signal. Using all epochs of maximum light calculated from the seasonal data in all filters, we derived the following ephemeris: T max = HJD (3) (2) E. (1) In Eq. (1), E denotes the number of elapsed cycles and the numbers in parentheses are the errors with the leading zeroes omitted. As we already pointed out in Paper I, our data suffer from systematic effects resulting in considerable power at low (<2 d 1 ) frequencies. This effect is even increased in the case of Oo 692 since the comparison star, Oo 843, is quite far from it. Therefore, we decided to decrease low frequency effects by removing linear trends from individual runs. Consequently, only long runs were used. The phase diagram of the resulting differential photomery of Oo 692 in all four filters is shown in Fig. 2, and the results of sine-curve fitting are given in Table 2. Although it is possible that Oo 692 does exhibit lowfrequency variations, we think that the signal we removed is mostly spurious, because of the systematic effects we mentioned, and the fact that no single low-frequency peak appears in the periodograms of a few separate subsets of our uncorrected data. The Fourier spectrum of the corrected data is shown in Fig. 3a. After prewhitening the data with the d 1 frequency, no other signals with amplitude exceeding the noise level ( 2 mmag) are seen (Fig. 3b). In the spectral window (Fig. 3c and d) both the daily and yearly aliases appear to be of similar height; many other peaks of intermediate frequencies are also seen. All these peaks, however, do not exceed 75% of the height of the main peak. This means that the ambiguity in deriving periods for Oo 692 and other short-period variables from our data is rather small.

3 J. Krzesiński et al.: -type pulsators in the central region of NGC 869 (h Persei) 57 Oo 692 U AMPLITUDE [mmag] Oo F R E Q U E N C Y [d ] Fig. 4. The same as in Fig. 3a and b, but for Oo 992. Prewhitening was performed with the period of d, corresponding to a frequency of d 1 (arrowed). I Oo 992 U PHASE Fig. 2. Differential UIlight curves of Oo 692 folded with the period of d. The ordinate ticks are separated by.1 mag. Phase. corresponds to HJD AMPLITUDE [mmag] AMPLITUDE [mmag] 1 5 Oo F R E Q U E N C Y [d ] a. b. c. d F R E Q U E N C Y [d ] Fig. 3. a Fourier frequency spectrum of the Oo 692 -filter nightly data corrected for linear trends. b The spectrum after prewhitening with the period of d. The ordinate scale in this panel is the same as in panel a. c Spectral window of the -filter observations of Oo 692 shitfed to frequency f = d 1 and normalized to the same amplitude as in the upmost figure. d A part of the spectral window around the frequency f showing the alias pattern including the yearly aliases PHASE Fig. 5. Differential U light curves of Oo 992 folded with the period of d. The ordinate ticks are separated by.1 mag. Phase. corresponds to HJD Oo 992 = D Schild (1965) gives the spectral type of Oo 992 as 1 n. The star is another β Cephei variable belonging to h Persei. The amplitude of the only period we have found above the noise level (Fig. 4) is extremely small (see Table 2 and Fig. 5). The I- filter data for this star were not used because of the large scatter. Despite the small amplitude we are sure that the variations seen in Fig. 5 are real because the same d 1 signal was found in a few separate subsets of our data. This time, the ephemeris for the time of maximum light was derived from the -filter data only: T max = HJD (8) (5) E. (2) The differential magnitudes, shown in the phase diagram (Fig. 5), were determined using five comparison stars (Oo 936, 963, 977, 98, and 14). ecause they are all relatively close to Oo 992, the residual standard deviation as well as the noise level in periodogram are considerably smaller than in the case of Oo 692. Like for Oo 692, we used only long runs corrected for linear trends.

4 58 J. Krzesiński et al.: -type pulsators in the central region of NGC 869 (h Persei) Table 2. The results of the sine curve fitting to the UI data of the three pulsating stars found in h Persei. N obs is the number of individual observations and RSD is the residual standard deviation. Fitted parameters have the same meaning as in Eq. (1) of Paper I with T = HJD Period P Semi-amplitude A Phase φ HJD of the time RSD Star N obs [days] Filter [mmag] [rad] of maximum light [mmag] Oo U 7.8 ± ± ± ± ± ± ± ± ± I 9.7 ± ± ± Oo U 1.8 ± ± ± ± ± ± ± ± ± Oo U 47.8 ± ± ± ± ± ± ± ± ± I 18.9 ± ± ± Oo 893 U This star appeared to be variable with a period of about 1.2 d. No spectral type is available for Oo 893, but its magnitude and colours are consistent with a mid- type star lying on the cluster main sequence. As can be seen in Fig. 6, the photometric amplitude in the U band is almost twice as large as in the three remaining ones. This feature is typical for all SP stars observed by Waelkens (1991). On the other hand, as is argued by Jerzykiewicz & Sterken (1993), it would be very difficult to explain such amplitude behaviour by means of variations in an ellipsoidal system. The period we found for Oo 893 is also characteristic for λ Eri variables. However, because of the small scatter in the light curve and the fact that our photometric Hα observations (Pigulski et al., in preparation) have not revealed any Hα emission in the star, this explanation seems to be improbable. We therefore conclude that, despite its evident monoperiodicity, Oo 893 is an SP variable. This is the first such a star found in this cluster. As can be seen in Fig. 7, after removing the main frequency (f = d 1 ) and its harmonic, 2f, there is no significant signal above the noise level (2 3 mmag). Ephemeris for the epoch of maximum light, derived from the and -filter data of Oo 893 is following: T max = HJD (3) (15) E. (3) PHASE Fig. 6. Differential UI light curves of Oo 893, folded with a period of d. The ordinate ticks are separated by.2 mag. Phase. corresponds to HJD Oo 893 I 4. Eclipsing binaries Three eclipsing binaries were discovered in the field. Their phase diagrams are shown in Fig. 8. The first one, Oo 121 = NS 779, was found to be probably variable by Oosterhoff (1937). Four well-defined minima were observed for this star. The corresponding times of minimum are given in Table 3. The most probable value of the orbital period is d. The primary minimum is only slightly deeper than the secondary and therefore there is a small possibility that the real orbital period is half the value given above. This, however, would imply a very shallow, i.e., practically unobservable secondary minimum and a very small mass ratio in the system. This is rather unlikely. Although barely visible in Fig. 8, the end of the egress was also observed on HJD in the U band (phase.76). There were, however, too few points falling into the eclipse to derive reliably the time of minimum for this night. The times of minimum light given in Table 3 define the following ephemeris: Min I = HJD (1) (1) E. (4)

5 J. Krzesiński et al.: -type pulsators in the central region of NGC 869 (h Persei) 59 AMPLITUDE [mmag] Oo F R E Q U E N C Y [d ] Fig. 7. The Fourier spectra of the differential -filter data of Oo 893. Upper panel: original data; middle panel: data prewhitened with frequency f = d 1 ; lower panel: data prewhitened with f and its first harmonic, 2f. The ordinate scale is the same for all three panels. Oo 121 Oo 1147 W PHASE Fig. 8. Phase diagrams for three eclipsing binaries discovered in h Persei. Upper panel: UI light curves of Oo 121 folded with period d; middle panel: light curve of Oo 1147 folded with period d; lower panel: light curve of W 49 folded with period d. For all stars phase. corresponds to HJD Ordinate ticks are separated by.2 mag. For Oo 1147, the other EA-type eclipsing binary, only one primary and one secondary minimum (see Table 3) both in and bands were observed. They are not sufficient to derive the period unambiguously. Assuming a circular orbit we found two possible orbital periods: ±.8 and ±.2 d. If the latter is real, a few points at the beginning of HJD fall into the secondary eclipse. Despite the lower quality of this night, a small decrease of brightness at that time is indeed observed. If small ellipticity is allowed, more periods, including U I Table 3. Heliocentric times of minimum light for three eclipsing variables in the central region of h Persei. T min (HJD 2 4 +) Filter Minimum Oo 121: ±.1 primary ±.2 secondary ±.6 primary ±.3 primary ±.6 I primary ±.3 secondary ±.3 secondary ±.4 I secondary Oo 1147: ±.5 primary ±.2 primary ±.7 secondary ±.4 secondary W 49: ±.3 primary ±.2 primary ±.2 primary ±.2 primary ±.7 primary ±.5 primary ±.4 primary some around 1.58 d, become possible. The eclipses are partial and the depths of the primary and secondary eclipses in the band amount to about.53 and.21 mag, respectively. Another eclipsing variable was detected only in the -filter data and it appeared to be a W UMa-type eclipsing binary with well-marked minima of uneven depth. The star has no Oosterhoff (1937) number. It is designated as star number 49 by Wildey (1964) and 475 by Moffat & ogt (1974). Taking into account the star s brightness and colours, it could be an early G-type star belonging to h Persei. The eclipse is probably partial. All times of minimum given in Table 3 define the following ephemeris: Min I = HJD (4) (1) E. (5) 5. Other variables Of remaining four variables, the most interesting is Oo 922 (D ) which appears to be a periodic one. The star was classified as.5n by Schild (1965). Periodogram analysis of the data of Oo 922 revealed a single peak at frequency 1.38 d 1. The period, derived from the most numerous -filter data, was found to be ±.12 d. Within the errors, the semi-amplitudes are almost the same in all three bands of the U system, amounting to 6.8 ±.6, 5.8 ±.3, and 6.1 ±.2 mmag for U,, and, respectively. In the 1997 data, however, the.725-day periodicity can barely be found; the changes seem to be rather aperiodic. From the characteristics of the star s variations, Oo 922 could be a λ Eritype variable. Fortunately, in late 1989 and early 199 the star was searched for the presence of the Hα emission by Goderya

6 51 J. Krzesiński et al.: -type pulsators in the central region of NGC 869 (h Persei) & Schmidt (1994). Their observations were carried out on two epochs separated by about three months. The first observation did not indicate the presence of emission, but the second one did. The Hα observations of h Persei were also carried out in iaĺków in the beginning of 1997, and no Hα emission was found in Oo 922 (Pigulski et al., in preparation). It seems that Oo 922 is a e star with emission which sometimes weakens to a level undetectable by photometric methods. It should be also remembered that λ Eri-type photometric variations are sometimes observed in n stars without emission (see alona 199). The brightest star in the field, the cluster supergiant Oo 157 (HD 14134, D , 52 Per), was classified as a 3 Ia star by Johnson & Morgan (1955) and as a 3 Iab star by Slettebak (1968). The star was suspected to be variable by Rufener & artholdi (1982) on the basis of their Geneva photometry. Its variability was later confirmed by Waelkens et al. (199). Oo 157 was also observed by the Hipparcos satellite. Although light changes are clearly visible both in the Hipparcos and our data, we failed to find any well-defined periodicity. The variations seem to be irregular with a range of about.15 mag in Hipparcos H p, and.6 mag in our data. In the periodogram of the data of Oo 986, the strongest peak appears at a frequency f 1 of (5) d 1. After prewhitening with this frequency, the next strongest peak occurs at f 2 =.62529(7) d 1, which is commensurable with f 1, i.e., f 1 3f 2. The cause of the variability is unknown. Percy (1972) flagged the variability of the 2 n star Oo 18 as uncertain and gave an upper limit of 5 mmag for its photometric -filter variations. Our observations confirm that the star is indeed variable. Analysis of all our -filter data yielded the periodicity of d. However, phasing with twice period ( d) results in a slightly smaller scatter. If the longer period is real, the most reliable explanation is that the star is ellipsoidal. Nevertheless, the possibility that Oo 18 is an SP, or even λ Eri variable similar to Oo 922, cannot be rejected. Assuming a 1.5-day period, the semi-amplitudes we derived were 9.9 ± 1.6, 6. ±.5, and 7.7 ±.3 for U,, and, respectively. Three other early- stars suspected to be variable by Percy (1972), namely Oo 963 (NS 776), Oo 978 (NS 777), and Oo 14 (NS 778), appeared to be constant. The detection thresholds for the full amplitude of periodic variations in for these three stars were 5, 6, and 6 mmag, respectively. Another of Percy s suspects, Oo 1161 (NS 781), was slightly outside the observed field. 6. U photometry 6.1. Transformations As was already mentioned in Sect. 2, our observations were carried out from two sites: Mt. Suhora (UI) and iaĺków ( I). 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. ecause the U 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). ecause of the lack of red photoelectric standards, we also applied the photographic photometry of some bright, well-isolated stars from Moffat & ogt (1974) giving them half the weight given to photoelectric measurements. The resulting transformation equations follow: = v.31(.35) (b v)+1.478(.6), (6) =.944(.34) (b v)+.334(.6), (7) U = 1.36(.41) (u b).461(.13), (8) 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.31,.44, and.59 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 iaĺ ków photometry with equations: = v.67(.12) (b v)+1.478(.1), (9) =.755(.8) (b v)+.335(.1). (1) This time, the lowercase letters stand for iaĺ ków instrumental magnitudes. The standard deviations are equal to.24 and.17 mag for Eqs. (9) and (1), respectively. Subsequently, for stars observed in both sites, transformed magnitudes and colour indices were combined together. Resulting standard magnitudes and colour indices for 184 stars as well as U 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 Here for reader s convenience we present only the data for 1 variable stars. The I-filter observations were not transformed to the standard Cousins I C system because there are no I C 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 = v +.76(.6) (v i) (.2) (11) =.683(.1) (v i)+.328(.3) (12) for iaĺków and = v.22(.2) (v i) (.1) (13) =.632(.11) (v i)+.342(.3) (14) for Suhora give us additional photometry for 74 faint stars. The standard deviations for Eqs. (11) (14) are, in succession, equal to.14,.24,.4, and.25 mag. This photometry is also included in Table 4. One may doubt the reality of the

7 J. Krzesiński et al.: -type pulsators in the central region of NGC 869 (h Persei) 511 Table 4. U 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 & ogt (1974) number, (6), magnitude, (7), colour index, (8), reference to photometry ( S means that the star was observed from Suhora only,, from iaĺków only, S+, both from Suhora and iaĺków. The photometry transformed with the use of I-filter measurements [Eqs. (11) (14)] are additionally flagged with I ), (9), U colour index, (1), right ascension (epoch 2.), (11), declination (epoch 2.), and (12), remarks. X Y Oo W M Ref U α 2. δ 2. Remarks (1) (2) (3) (4) (5) (6) (7) (8) (9) (1) (11) (12) S S S D , HD S S D S S D S S D colour indices transformed from v i 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. ecause the mean instrumental magnitudes were calculated from all search frames, the internal accuracy of our U photometry is very good it is better than 1 mmag for brightest stars and reaches about.2 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 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 v i indices is shown with plus signs in the CM diagram. As in χ Persei, both the h Persei β 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 λ Eri star Oo 922. The magnitudes of the β Cephei stars (9.35 for Oo 692, and 9.9 for Oo 992) are in the same range as in χ 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 U colour index for this star is equal to.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 U 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., E(U )/E( )= E( ). (15) The best agreement was obtained for E( ) =.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 E( ) for h Persei differs from.52 mag by no more than.1 mag. The average value of E( ) is in general agreement with previous determinations (see Tapia et al and 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.5 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 β Cephei stars we discovered. Therefore, we could not convert Strömgren indices to log T eff 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 Comparison with previous work We compared our U photometry with previous photoelectric (Johnson & Morgan 1955; Schild 1965; Tapia et al. 1984)

8 512 J. Krzesiński et al.: -type pulsators in the central region of NGC 869 (h Persei) ( - ) (U-) -.6 a. b ( - ) Fig. 9. a The colour-magnitude diagram for the observed field in h Persei. The two β Cephei stars are denoted by open, and the other stars by filled circles. ariables are enclosed in large open squares. Stars for which photometry was transformed from the instrumental I measurements are shown with plus signs. b Colour-colour diagram for stars which U photometry is available. Our photometry is shown with the same symbols as in the left diagram. Crosses are points in which the U 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, E( ) =.52 mag (long solid line), and E( ) =.47 and.57 mag (short solid lines) are also plotted. 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. oth the mean differences and SD are expressed in mag. Photometry SD N ( ) SD N (U ) SD N Johnson & Morgan (1956).3 a.36 a a.3 a Schild (1965) Tapia et al. (1984) Wildey (1964) Moffat & ogt (1974) 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 ( ). and photographic (Wildey 1964; Moffat & ogt 1974) U studies. The result of these comparisons is shown in Fig. 1, and the mean differences are given in Table 5. Our photometry agrees quite well with all photoelectric studies. Out of the two photographic U data sets, the one of Moffat & ogt (1974) is much better. The differences between our and Wildey s (1964) 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 & ogt (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. 7. Discussion and conclusions The discovery of β Cephei pulsations in Oo 692 and Oo 992, the two early -type stars located in the central region of h Persei, together with our findings described in Paper I mean, that the double cluster h and χ Persei is indeed rich in these kind of variables. Since our up-to-date observations were confined to the central regions of the double cluster, more stars of this type may be discovered. Regarding the richness of h and χ Persei, we estimated that the clusters should contain about 1 15 β Cephei stars. Candidates have already been selected and observations of other fields in both clusters are now under way.

9 J. Krzesiński et al.: -type pulsators in the central region of NGC 869 (h Persei) 513 [mag] (-) [mag] (U-) [mag] [mag] Fig. 1. Comparison of different U photometries. All the differences are in the sense our minus Johnson & Morgan (1955, filled circles), Schild (1965, open circles), Tapia et al. (1984, open squares), Wildey (1964, plus signs), and Moffat & ogt (1974, dots). Another interesting discovery in h Persei is finding of SPtype variations in Oo 893. Although the star is monoperiodic within the detection limit of our observations, the distinct wavelength dependence of the amplitude of variations is a strong argument in favour of an SP classification. Many monoperiodic stars found in the Hipparcos data (Waelkens et al. 1998) were also classified as SPs. It seems, therefore, that multiperiodicity cannot be longer regarded as a necessary condition for classifying a star as an SP. Consequently, some of the monoperiodic variables found by alona and co-workers in three southern young open clusters (alona 1994; alona & Koen 1994; alona & Laney 1994) could be, in fact, SP stars. Acknowledgements. This research was supported by the KN grant No. 2 P3D We would like to thank our friend Liza van Zyl for critical reading the manuscript. References alona L.A., 199, MNRAS 245, 92 alona L.A., 1994, MNRAS 267, 16 alona L.A., Koen C., 1994, MNRAS 267, 171 alona L.A., Laney C.D., 1994, MNRAS 276, 627 Caldwell J.A.R., Cousins A.W.J., Ahlers C.C., van Wamelen P., Maritz E.J., 1993, SAAO Circ. 15,1 Franco M.L., Magazzù A., Stalio R., 1985, A&A 147, 191 Goderya S.N., Schmidt E.G., 1994, ApJ 426, 159 Jerzykiewicz M., Sterken C., 1993, MNRAS 26, 826 Johnson H.L., Morgan W.W., 1955, ApJ 122, 429 Krzesiński J., 1995, PASPC 83, 299 Krzesiński J., Pigulski A., 1997, A&A 325, 987 (Paper I) Moffat A.F.J., ogt N., 1974, eröffentlichungen des Astronomischen Instituts der Ruhr-Universität ochum Nr. 2 Natali F., Natali G., Pompei E., Pedichini F., 1994, A&A 289, 756 Oosterhoff P.T., 1937, Ann. Sterrewacht Leiden 17, 1 Pandey A.H., hatt.c., Mahra H.S., Sagar R., 1989, MNRAS 236, 263 Percy J.R., 1972, PASP 84, 42 Rufener F., artholdi P., 1982, A&AS 48, 53 Schild R.E., 1965, ApJ 141, 979 Slettebak A., 1968, ApJ 154, 933 Stetson.P., 1987, PASP 99, 191 Tapia M., Roth M., Costero R., Navarro S., 1984, Rev. Mex. Astron. Astrofis. 9, 65 Turner D.G., 1989, AJ 98, 23 Waelkens C., 1991, A&A 246, 453 Waelkens C., Lampens P., Heynderickx D., et al., 199, A&AS 83, 11 Waelkens C., Aerts C., Kestens E., Grenon M., Eyer L., 1998, A&A 33, 215 Wildey R.L., 1964, ApJS 8, 439