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1.
A first period study of the eclipsing binary XY Ceti is presented. A new period (P=2d.7807135), based on all available times of minima, is given. Period changes in different portions of the O–C diagram, with a new period, have been estimated. The total change in period (P/P) ranges from 1.1×10–5 d to 1.2×10–4 d, thus, P ranges from 3.1×10–5 d to 3.3×10–4 d. The O–C diagram suggests that the trend of the period has changed around the year 1959. Two portions of increasing and decreasing trends also reveal that the period changes (P/P) of the order of 10–5 d are present, which are appreciably large.  相似文献   

2.
A new period (P=8d.439422) of the eclipsing binary system EI Cephei has been given, which is based on all available times of minima. Periods using Strohmeier's (1958) epoch have also been presented for the observations given by other investigators. Period based on only photoelectric minima comes out to be 8d.439336, which is lesser than the earlier periods given in the literature. O-C diagrams of EI Cephei have been presented for the first time, and period variations have been estimated in different portions of the O-C diagram (Figure 2) of EI Cephei. Strong period changes have occurred around the years 1959 and 1965. The total change in period (P/P) ranges from 6.7×10–5 to 4.3×10–4. Thus, P of the order of 10–3 d are present, which fact suggests that strong period variations are present in EI Cephei. However, periods given by various investigators show no systematic trend of period variations. The existence of a third body in the system could not be confirmed.  相似文献   

3.
The period study of the eclipsing binary DF Hya, based on up-to-date minima has been presented. The least-squares method has been applied to obtain a new period, which comes out to be 0 . d 3306017. Period changes are found around the years 1949, 1974, and 1982. Appreciable period changes are apparent around 1949 and 1974, the strongest being around 1974. The period changes (P) range from 0.46×10–6 d to 0.46×10–5 d, the average being 1.89×10–6 d. Such period changes are usual for a contact system, like DF Hya. Our results do not show increase in the time interval 1959–1985 as suggested by Zhanget al. (1989).  相似文献   

4.
A new period (P=1d.7175405) of the eclipsing binary system SZ Arietis has been presented. Period changes in different portions of the O-C diagram, with new period, have been estimated. The total change in period (P) ranges from 3.64×10–5 d to 4.24×10–4 d, which is appreciably large. However, leaving the unusual value, the average period change comes out to be of the order of 6×10–5 d. The period changes around the years 1903, 1943, and 1977 are apparent in the O-C diagrams. A sinusoidal variation is also visible in the O-C diagrams which indicates that SZ Ari may be a three-body system, having a period of nearly 66 years.  相似文献   

5.
The new period (P=0 . d 461700) of the eclipsing binary system DX Aqr has been presented, which is based on available times of minima. O–C diagram of DX Aqr has been presented for the first time, and the period variations present in the system have been analysed. In all five period increases and five period decreases are nothed, and four period increases and five period decreases have been discussed. The strongest period increase occurs between 1975 and 1976. The total period change in different portions of the O–C diagram ranges from 1.40×10–4 d to 3.61×10–6 d. Appreciable period fluctuations have been noted to have occurred in the time intervals, 1964–1965 and 1974–1975.  相似文献   

6.
The period study of the eclipsing binary system GH Pegasi has been presented for the first time. A new period (P=2d.556135) of GH Peg, based on all available times of minima, has been given. O–C diagrams of the system have also been presented for the first time, and the period changes present in the system have been analysed. The period shows changes around the year 1972 and 1981. The total period change in different portions of the O–C diagram, based on the corrected period, ranges from 5.2×10–6 d to 7.0×10–5 d. The photoelectric minima show sufficiently large scatter in the system.  相似文献   

7.
A detailed period study of the eclipsing binary system AW UMa is presented. A new period (P=0d.4387317) is given. Period changes in different portions of the O-C diagram, based on new period, have been estimated. The total change in period (P) ranges from 2.2×10–7 to 2.8×10–6 d, which is normal for AW UMa systems. Two distinct linear trends of period are visible in the O-C diagram. A third trend shows a new change in the period behaviour, which is yet to be confined.  相似文献   

8.
Detailed period study of the eclipsing binary ST Per is presented. A new period (P=2d.648339) is given. Period changes in different portions of the O-C diagram with a new period have been estimated. The total changes in period (P) ranges from 2.17×10–5d to 2.64×10–4d which is appreciably large. Sufficient number of minima in the time interval 1934 to 1985 for this system are available. Distinct increasing and decreasing trends are evident, the change in the tendency appears to have occurred around 1947. Sinusoidal variation is seen between cycles 7000–10000, which indicates that ST Per is a three-body system, the period of the third body being about 22 years. However, the sinusoidal variation is not perfectly symmetric in shape, therefore, it is suspected that ST Per is a four-(or multi-) body system.  相似文献   

9.
On the basis of observational data for the absolute R and relative R/R amplitudes of variations in radius of galactic classical cepheids (55 stars from Balona and Stobie (1979) and 30 stars from Sollazzoet al. (1981)), four kinds of empirical linear relations are obtained: log(P V)–logR, logP–logR, log(P V)–log(R/R), and logP–log(R/R);P, R, and V are the pulsation periods, the mean stellar radii, and the amplitudes of light variations, respectively. Three groups of stars are considered: short-period cepheids (SPC)-with logP1.1; long-period cepheids (LPC)-with logP>1.1; and s-cepheids (sC). Both the R values and the R/R values increase withP andP V, for a given group of variables. A comparison is performed with our results obtained from data in other sources (Kurochkin, 1966; Gieren, 1982; etc.). The investigated relations can be applied for determining R and R/R of galactic classical cepheids, by using their observedP and V. All studied galactic classical cepheids have R/R<0.35, R<10R for SPC and 10R <R60R for LPC. The sC have smaller R and R/R values than other classical cepheids, at the same periods (the difference is about 2 times for R and 1.4–2.8 times for R/R); the studied sC have R/R in the range 0.025–0.075 and R in the range 1–3R (only Y Oph has R8R ).  相似文献   

10.
A new period (P=3d.687664) of the eclipsing binary system IZ Persei is given, based on 16 observed times of the minima. O–C diagrams of IZ Per have been presented for the first time, and the period variations have been estimated in different portions of the O–C diagram. Significant period changes do not appear to have occurred in IZ Per. The O–C diagrams suggest that the period of the system is continuously increasing at a rate of 25s yr–1. Period variations of the order of 10–5 d appear to have occurred around the years 1969, 1972, and 1978. The period increases are stronger than the period decreases; but these are yet to be confirmed. The overall picture of IZ Per suggests that strong period changes are not present in the system; however, slow increase of period is apparent in IZ Per. The total period change (3×10–6 d) till the last epoch is in agreement with the newly derived period of IZ Per.  相似文献   

11.
New period (P=3d.067182) of the eclipsing binary system XX Cassiopeiae has been given, which is based on thirty-nine minima. O–C diagram of XX Cas has been presented for the first time, and the period variations present in the system have been analysed. In all seven period increases and four period decreases are noted in the time-interval 1922 to 1983. Strongest period change occurs in the time-interval 1929 to 1933. The total period change in different portions of the O–C diagram ranges from 6.1×10–7 d to 6.3×10–4 d. Appreciable period fluctuations are noted in the time-interval 1922 to 1936. The possibility of the presence of a third body has been suspected.  相似文献   

12.
With the aid of the spectra taken in the years 1959–1968, a physical analysis of the atmosphere of P Cygni has been carried out and the motions of the atmosphere have been studied. The variations of radial velocities, the velocity progressions of Balmer and Hei lines, the high rate of mass loss (2×10–5 M yr–1), the features of the observed line profiles, especially that of H-K lines of Caii andD 1-D 2 lines of Nai confirm the conclusion of Van Blerkom (1978), concerning the assumption of an accelerating atmosphere for P Cygni. The electron density variation with the radius seems to ben e r –5/2, with an average value of 7×1011cm–3 at the lower boundary of the atmosphere.In order to explain the two absorption components of observed lines, an atmospheric model based on the assumption of three envelopes, two of which accelerate gradually with two different velocity laws (up to 11.2r c ), and the third of which accelerates rapidly with a standard velocity law (beyond 11.2r c ) has been developed. From this model and the observed profiles, the geometrical thicknesses of the line-forming regions of H, H, H, and H are derived.The observations were obtained at Haute Provence Observatory (CNRS).  相似文献   

13.
For nine published high-resolution QSO spectra a correlation analysis of their L forest lines has been performed. The two-point correlation functions show some quasi-periodic structure of magnitude ||0.3. Their characteristic separation along the line-of-sight amounts to s 0=3×10–3 or to s 0=5×10–3 for =1 and 0.2, respectively. Especially the distribution of nearest neighbouring line positions in two close QSO pairs allows for the interpretation that the absorption clouds lie in sheet-like structures as predicted by the pancake theory. The correlation data contain some hints on metal absorbers within the forest of unidentified lines.  相似文献   

14.
A comprehensive period study of the times of minima observed from 1881 to 1985 has been performed. Previous interpretations of the O–C diagram based on light-time effect are confirmed. The light-time orbit of U Oph has been revised using a differential corrections procedure. We get an eccentric orbit withe=0.22±0.06,P=38.7±0.2 yr, and a mass function . In addition, our analysis has revealed short-period apsidal motion (U/P=4515±75) in a slightly eccentric close orbit (e=0.0031±0.0003), allowing a reliable determination of the density concentration coefficient,k 2=0.0059±4. A comparison with stellar evolutionary models calculated by Jeffery (1984) yields the helium contentY=0.28±0.05 and an age of 3×107 yr for the components of U Oph.  相似文献   

15.
Wheatland  M.S.  Litvinenko  Y.E. 《Solar physics》2002,211(1-2):255-274
The observed distribution of waiting times t between X-ray solar flares of greater than C1 class listed in the Geostationary Operational Environmental Satellite (GOES) catalog exhibits a power-law tail (t) for large waiting times (t>10hours). It is shown that the power-law index varies with the solar cycle. For the minimum phase of the cycle the index is =–1.4±0.1, and for the maximum phase of the cycle the index is –3.2±0.2. For all years 1975–2001, the index is –2.2±0.1. We present a simple theory to account for the observed waiting-time distributions in terms of a Poisson process with a time-varying rate (t). A common approximation of slow variation of the rate with respect to a waiting time is examined, and found to be valid for the GOES catalog events. Subject to this approximation the observed waiting-time distribution is determined by f(), the time distribution of the rate . If f() has a power-law form for low rates, the waiting time-distribution is predicted to have a power-law tail (t)–(3+) (>–3). Distributions f() are constructed from the GOES data. For the entire catalog a power-law index =–0.9±0.1 is found in the time distribution of rates for low rates (<0.1hours –1). For the maximum and minimum phases power-law indices =–0.1±0.5 and =–1.7±0.2, respectively, are observed. Hence, the Poisson theory together with the observed time distributions of the rate predict power-law tails in the waiting-time distributions with indices –2.2±0.1 (1975–2001), –2.9±0.5 (maximum phase) and –1.3±0.2 (minimum phase), consistent with the observations. These results suggest that the flaring rate varies in an intrinsically different way at solar maximum by comparison with solar minimum. The implications of these results for a recent model for flare statistics (Craig, 2001) and more generally for our understanding of the flare process are discussed.  相似文献   

16.
The O–C diagram of the eclipsing binary GG Cassiopeiae has been presented for the first time, and the period changes present in the system have been analysed. In all three period changes are noted. The strongest period change has been found to occur in the time-interval 1942 to 1966. The total period change in different portions of the O–C diagram ranges from 7.1×10–7 d to 2.0×10–5 d. The stronger period changes appear to have occurred after 1942; prior to it, the system has shown a negligible period change. The overall picture of the O–C diagram suggests that the O–C values of the system GG Cas are negative after 1942. The presence of a third body does not appear probable. The period fluctuations are also appreciable. A new period (P=3 . d 758733) has been presented.  相似文献   

17.
P. Maltby 《Solar physics》1977,55(2):335-346
The effects of the magnetic field as well as the velocity field on sunspot equilibrium are discussed. The gas pressure difference, P, between a spot and the environments in the same horizontal layer is primarily determined by the magnetic field. Using recent model atmospheres we find that P shows a maximum value, P max, at a depth of 650 ± 150 km below the photosphere. The value of P max suggests that the curvature of the field lines is important for the equilibrium.It appears that, at an optical depth of unity in the umbra, the density has a value close to that of the environment at the same geometric depth (see Figure 4). If such is the case the expression for the umbra temperature (Equation (15)) may be considerably simplified (Equations (17) and (18)) and compared with observations.  相似文献   

18.
A detailed investigation of the evolution of low-mass binaries is performed for the case when the secondary fills its Roche lobe at the stage of core hydrogen exhaustion. The obtained results are compared with observational data for ultra-short periodic X-ray systems MXB 1820-30 and MXB 1916-05. In the frame of the proposed evolutionary scenario it is possible to obtain for MXB 1820-30 its periodP=11.4 min twice (see Figure 2). In the first case the parameters of the system are:M 2 0.13–0.15M ,X0.05–0.13, |P/P| (3.6–6.2) } 10–7 yr–1, M2 (4.1–9.6) } 10–9 M yr–1, for the second:M 2 0.08–0.09M ,X= 0, |P/P| (1.3–1.5) } 10–7 yr–1, M2 (1.4–1.8) } 10–8 M yr–1. It is suggested that MXB 1916-05 is the progenitor of the system MXB 1820-30 (M 2 = 0.1M,X 0.221,M 2 1.8 × 10–10 M yr–1).  相似文献   

19.
The period variations of TV Cassiopeiae between 1901 and 1977 are discussed in the light of the period change model proposed by Biermann and Hall. During each period decrease 4.0×10–6 M of mass is transferred from the secondary star to the primary. The average observable mass transfer rate is found to be 4.3×10–7 M yr–1. This average rate corresponds to the thermal time-scale of the mass-losing star.  相似文献   

20.
E. Kirsch 《Solar physics》1973,28(1):233-246
Solar neutron emission during large flares is investigated by using neutron monitor data from the mountain stations Chacaltaya (Bolivia), Mina Aguilar (Argentine), Pic-du-Midi (France) and Jungfraujoch (Switzerland). Registrations from such days on which large flares appeared around the local noon time of the monitor station are superimposed with the time of the optical flare as reference point.No positive evidence for a solar neutron emission was found with this method, However, by using an extrapolation of the neutron transport functions given by Alsmiller and Boughner a rough estimation of mean upper limits for the solar neutron flux is possible. The flux limits are compared with Lingenfelter's model calculations.From the Chacaltaya measurements it follows: N 02.8 × 10–3 N cm–2 s–1 per proton flare, E > 50 MeV, if P0 = 125 MV N 01.4 × 10–2 N cm–2 s–1 per proton flare, E > 50 MeV, if P 0 = 60 MV and from Pic-du-Midi measurements: N 06.7 × 10–3 N cm–2 s–1 per proton flare, E > 50 MeV, if P 0 = 125 MV N 04 × 10–2 N cm–2 s–1 per proton flare, E > 50 MeV, if P 0 = 60 MV P 0 = characteristic rigidity of the producing proton spectrum on the Sun.The flux limits estimated for some special proton flares are consistent with Lingenfelter's predictions for the acceleration phase but are too small for the slowing down phase. Therefore it is believed that Lingenfelter's assumption of isotropic proton emission from the flare region is not fulfilled.  相似文献   

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