3. Data analysis

3.1. Memberships

Photometric diagrams of observed data are shown in Figs. 4 (click here) to 6 (click here). In particular, the two-color diagram, reveals a cluster sequence clearly outlined up to but merging into field stars for larger values of B-V. The slight spread shown by the bluest stars in this diagram could be due to differential reddening as photometric errors or contamination by neighbors can be discarded. Also in this diagram, many heavily reddened blue stars are located above the reddening line suggesting, in principle, the presence of absorbing material behind the cluster. However, since some of them have large U-B errors (> 0.08 mag) their settlements in this diagram are definitely dubious and no firm conclusion can be drawn.

After inspecting all the photometric diagrams, membership assessment could be done realistically for stars down to mag. By comparison, mainly in the UBV diagrams, a total of 26 likely members were found down to 14 along with several probable members contained in the range . All likely members have single reddening solution in the U-B,B-V diagram and, by assuming that they are all of luminosity class V, reddening-free UBV colors were obtained by a standard procedure (Vázquez et al. 1994) where individual determinations of color excesses yielded mean reddenings and respectively. Probable members found for and V < 16 mag were all de-reddened with the mean color excesses derived from likely members.

The brightest stars of the cluster, #3, 5, 6 and 8, show changing positions in the photometric diagrams: they are all shifted to the right side in the diagram while in the , in addition to them, we found that #9 and 14 show an extra-displacement too. Besides, in the array, where there is no V-I index of star #3, we notice that stars #5 and 6 are also at the right side of the main sequence band. This fact could be explained by a variety of causes such as weak emission of B-type stars, unresolved companions, binarity, and, finally, Ap stars, any of them causing the stars to move from their positions in the main sequence band. Another stars, as #18, 28, 36 and 37 appear, however, subluminous (chiefly in the diagram) whereas stars #9, 14, 26, 39 and 53 are less affected by reddening.

Considering only the bright stars, the reddening increases from east to west and from south to north across the cluster following the noticeable dust distribution in Fig. 1 (click here) although it never surpasses a total variation of 0.30 mag.

Nonetheless, looking at Fig. 6 (click here)b another peculiarity emerges: most of likely members are placed between R-values from 3.1 to even larger than 3.6. Are these anomalous R-values produced by the interstellar material in the cluster environment?

When we investigated the neighbor cluster Hogg 16 (Vázquez & Feinstein 1991b), situated at of Cr 272, we found a normal value R=3.0 and a similar mean reddening . Therefore, it is not simple to think of a mechanism able to change suddenly the interstellar matter properties in a so small and trivial region. Indeed, spectroscopy and polarimetry of the bright members could help us to explain their peculiar R-values and red displacements.

Meanwhile, we will adopt a normal value R=3.1 to produce reddening-free magnitudes in the diagram of Fig. 7 (click here) as .

3.2. Distance and age

The superposition of the ZAMS (Schmidt-Kaler 1982) in the corrected diagram of Fig. 7 (click here) gives a distance modulus that situates the cluster at pc from the Sun and close to the outer border of the Becker's inner arm-II, as suggested by Fenkart et al. (1977). The latter authors found a cluster distance of d=2800 pc, the difference with ours being produced, mainly, by the uncertainty of the ZAMS fitting in the scattered cluster sequence obtained by them. If instead of R=3.1 we had used R=3.4 (an average of Fig. 6 (click here)b), the distance of Cr 272 would have been d=3200 and our conclusions would not change at all.

With the cluster distance modulus we obtained the mag of the likely members listed in Table 3 (click here) together with their intrinsic colors and excesses. We also show in the table the "equivalent" spectral types based upon the Schmidt-Kaler (1982) calibration. The cluster earliest spectral type could be situated between "b1" and "b2", later than "o5" proposed by Fenkart et al. (1977).

   

F Phot. ST Rem

5 15 -0.27 0.56 -0.96 0.42 b1.5v 8.95 -2.88 (1)

6 10 -0.23 0.46 -0.84 0.34 b2v 9.60 -2.24 (2)

8 70 -0.22 0.51 -0.78 0.38 b2.5v 9.79 -2.05

9 50 -0.17 0.37 -0.59 0.28 b5v 10.27 -1.57

14 45 -0.17 0.38 -0.56 0.28 b5v 10.64 -1.20

18 103 -0.24 0.48 -0.87 0.36 b3v 10.52 -1.32 (3)

23 11 -0.19 0.48 -0.67 0.36 b3.5v 10.89 -0.95

25 -0.19 0.47 -0.64 0.35 b5v 10.97 -0.87

26 -0.13 0.33 -0.43 0.24 b7.5v 11.55 -0.29

28 -0.20 0.45 -0.69 0.34 b4v 11.26 -0.58 (3)

32 -0.10 0.43 -0.28 0.32 b8.5v 11.73 -0.11

33 -0.15 0.46 -0.52 0.34 b6.5v 11.65 -0.19

36 101 -0.17 0.47 -0.57 0.35 b6v 11.79 -0.05 (3)

37 100 -0.17 0.45 -0.56 0.34 b6v 11.86 0.02 (3)

39 35 -0.07 0.35 -0.18 0.26 b9/a0v 12.25 0.41

40 -0.12 0.46 -0.37 0.34 b8/b9v 11.98 0.14

41 -0.07 0.45 -0.18 0.34 b9/a0v 12.03 0.19

43 -0.14 0.57 -0.46 0.43 b7/b8v 11.76 -0.08

46 -0.11 0.43 -0.32 0.32 b9/a0v 12.26 0.42

48 65 -0.10 0.42 -0.29 0.31 b8/b9v 12.37 0.53

51 -0.08 0.48 -0.22 0.36 b9v 12.23 0.39

52 -0.11 0.40 -0.34 0.29 b8.5v 12.50 0.66

53 43 -0.05 0.38 -0.12 0.28 b9.5v 12.60 0.76 (2)

57 0.01 0.46 0.07 0.34 a0v 12.51 0.67

58 97 -0.11 0.46 -0.33 0.34 b8.5v 12.56 0.72

Note: The first column contains the star identification numbers used in this work and the second one contains the numbers from Fenkart et al. (1977). The spectral types shown in Col. 7 were derived from a combination of intrinsic indices and the resulting adopting the calibrations given in Schmidt-Kaler (1982).

Rem: (1), not measure available in Fenkart et al.; (2), non member according to Fenkart et al.; (3), slightly underluminous according to their intrinsic color indices.

By fitting the isochrones of Maeder & Meynet (1988) evolutionary models to the cluster upper sequence of Fig. 7 (click here), we found a cluster age of 15.8 Myr. As this fitting is somewhat uncertain, we inspected other possibilities such as the bluest color at the turnoff point (in ) which gives an age of 10 Myr according to the calibration of Meynet et al. (1993), on a side. Also, by interpolation in Maeder & Meynet models, the actual mass of the most luminous star, #3, is corresponding to an age of 14.3 Myr. Nevertheless, star #5, situated close to the ZAMS, has a mass of and its age is then 12.5 Myr. Therefore, we adopt Myr as a reasonable estimate of the cluster age.

Among the stars observed by Vázquez & Feinstein (1991b) in Hogg 16, four of them lie in the field of Cr 272. In particular, stars #53 and 71 in that paper (#6 and 18 in the present work, respectively) are, indeed, members of Cr 272 and not of Hogg 16.

The angular distance between Cr 272 and Hogg 16 is only ( pc) and Hogg 16 is, in addition, yr old, as seen in Fig. 7 (click here). Therefore, considering that, at level these clusters are at a same distance from the Sun and have similar reddenings, they could represent an example of sequential formation (Elmegreen & Lada 1977) where the bright stars of Hogg 16 triggered the star formation in Cr 272. Recently, Subramaniam et al. (1995) proposed that about 8% of the known galactic clusters could be members of binary systems. In this hypothetical case, we mention that Cr 272 and Hogg 16 have linear separation and age difference of the order of typical pairs listed by those authors.

As already noticed by Fenkart et al. (1977) no evolved stars are found in the cluster. Looking for them in the periphery of Cr 272, we found the star HD 117399, a Cep V659 Cen variable with spectral type F6/7Ib and period (Houk & Cowley 1975), situated at to the southwest. This star has no chance to be a cluster member because Evans (1992) and also Fernie et al. (1995) indicate that it has and a companion star of spectral type B6. Therefore, it does not fit into the scheme of Cr 272 as seen in Fig. 7 (click here).