Direct detection of the secondary spectrum of the interacting eclipsing binary TX UMa

Eclipsing binary TX UMa was observed with the D.A.O. high-dispersion spectrographs in 1969–1970, with emphasis on the detailed coverage of the primary minimum. One spectrum was taken exclusively within totality, thus exhibiting an uncontaminated spectrum of the secondary component. This leads to spectral reclassification of the secondary (F6 IV). The narrowing of the line profile of the H-line in totality is interpreted in terms of synchronous rotation of the secondary (v sini80 km s^–1) while the primary rotates faster (v sini130 km s^–1) than synchronously (v sini50 km s^–1). Although the secondary does not fill in its Roche lobe fully, the system exhibits pronounced indications of rather strong physical interaction. This is now supported also by the profound changes of the line profiles of the H-line with phase.     

Time Evolution of low-Frequency Periodicities in Cosmic ray Intensity

The long-time series of daily means of cosmic-ray intensity observed by four neutron monitors at different cutoff rigidities (Calgary, Climax, Lomnický tít and Huancayo/Haleakala) were analyzed by means of the wavelet transform method in the period range 60 to 1000 days. The contributions of the time evolution of three quasi-periodic cosmic-ray signals (150 d, 1.3 yr and 1.7 yr) to the global one are obtained. While the 1.7-yr quasi-periodicity, the most remarkable one in the studied interval, strongly contributes to the cosmic ray intensity profile of solar cycle 21 (particularly in 1982), the 1.3-yr one, which is better correlated with the same periodicity of the interplanetary magnetic field strength, is present as a characteristic feature for the decreasing phases of the cycles 20 and 22. Transitions between these quasi-periodicities are seen in the wavelet power spectra plots. Obtained results support the claimed difference in the solar activity evolution during odd and even solar activity cycles.     

On the spectrum of relativistic electrons accelerated in cosmic-ray sources

The data on the spectrum of the cosmic-ray electron component near the earth, on the radio-spectra of radio-galaxies, quasars and the Crab Nebula, as well as the data pertaining to the X-ray spectrum of the cosmic background, all agree that the sources of cosmic-ray electrons (such as supernovae and galactic nuclei) inject particles characterized by a power spectrumN(E)=KE ^–0, with 01.5–2.5. A mechanism is known in which the source emits a proton-nuclear component of cosmic rays with a spectrumN _n (E)=K _n E ^– _n, _n = + 2, =w cr/(w–w cr), wherew cr is the cosmic-ray energy density in the source, andw=w cr+w n+w _turb, the total energy density. We obtain =2.5 in agreement with observations on the natural assumption that =0.5. Within the framework of the same model with some additional assumptions, the electrons in the source, as well as those ejected by the source, are shown to have a power-spectrum characterized with _{0 n} = + 2. Thus the model discussed gives an adequate spectrum for both the proton-nuclear and the electron components of cosmic rays.     

Cosmic gamma-ray burst from intergalactic relativistic dust grains

Charged dust grains of radiia3×10^–63×10^–5 cm may acquire relativistic energy (>10¹⁸ eV) in the intergalactic medium. In order to attain relativistic energy, dust grains have to move in and out (scattering) of the magnetic field of the medium. A relativistic grain of radiusa10^–5 cm with Lorentz factor 10³ approaching the Earth will break up either due to electrostatic charge or due to sputtering about 150100 km, and may scatter solar photons via a fluorescence process. Dust grains may also melt into droplets in the solar vicinity and may contribute towards observed gamma-ray bursts.     

Statistical survey of planetary nebulae: Distances,masses, and distribution in the galaxy

On the basis of empirical (D)-dependency at the frequency of 5 GHz constructed using 15 planetary nebulae with the independently measured distances (10^–171×10^–20 W m^–2 Hz^–1 ster^–1), we evaluated distances of 335 objects. Independent evidence of the correctness of the accepted scale are given. Then(D)-dependency is constructed and it is shown that atD<0.08 pc the mean electron density is higher than the one determined by the Seaton method. We showed that the filling factor diminishes with the increase of the PN diameter (1 atD0.08 pc and 0.2 atD0.4 pc). the ionized mass of 33 PNs is determined. With the diameter increase the ionized mass grows and atD0.4 pc reaches the valueM0.07M . We used the new distance scale when investigating the space distribution of PNs. The mean scale height =130±15 pc and the mean gradient of the change of surface densitym=0.37, which allowed us to estimate the total number of nebulae in the GalaxyN4×10⁴. We divided the PNs according to their velocities (withV _LSR>35 km s^–1 andV _LSR<35 km s^–1) and permitted us to confirm that the PN belong to different sub-systems of the Galaxy. The estimated local formation rate of PNs [=(4.6±2.2)×10^–12 pc^–3 yr^–1] is a little higher than the one of the white dwarfs. That can be explained by a large number of PNs having binary cores, which used in our sample. The statistical estimation of PN expansion velocity showed thatV _ex increases from 5–7 km s^–1 (atD0.03 pc) to 40–50 km s^–1 (atD0.8 pc).     

Space-time profile variations of some spectral lines and their influence upon doppler velocity measurements

A technique for high-sensitivity measurements of spectral line profile fluctuations is suggested. Observations with spectral lines most commonly used to study the oscillations have been carried out. It is found that 5-min and 3-min fluctuations of Fei 5123, 5250, 5434 and NaDi 5896 line profiles are able to produce signals equivalent to line-of-sight velocities of 1–5 m s^–1 at a spatial resolution of 5 and 10–35 m s^–1 at 1.5 × 4 resolution. Such observations permit a better understanding of the particular physical factors responsible for the oscillations of line-of-sight velocity signals and the magnetic field which are the subject of study of helioseismology.     

Correlation of He II lyman alpha with He I 10830 Å, and with chromospheric and EUV coronal emission

This paper describes the results of comparing SERTS-3 images obtained in the transition region line of Heii 304 with chromospheric Hei 10830 absorption, with strong coronal lines of Mgix 368 , Fexv 284 and 417 , and Fexvi 335 and 31 , with H, with Caii 8542 , and with magnetograms in Fei 8688H. All of the images are illustrated, and the image reconstruction techniques used are described and evaluated. The high correlation of the Heii 304 and Hei 10830 images, originally found by Harvey and Sheeley (1977), is confirmed and is put on a quantitative basis. We find that the supergranulation network has greater contrast, and that filaments appear darker, in 10830 than in 304 . In active regions, the 304 line follows more closely the behavior of H and Caii 8542 than the 10830 line.     

Some limits on cosmic-ray heating ofHI clouds by magnetic stars

Estimates are made of the low-energy, cosmic-ray power generated by the rotational braking of magnetic Ap stars through interaction with the interstellar medium. The resulting cosmic-ray power density (1.4×10^–29 J s^–1 m^–3) in the galactic plane is at least three orders too small to maintain the observed kinetic temperatures (100 K) ofHi clouds. About 0.4% ofHi clouds contain an Ap star; in such cases 10²⁶ J s^–1 may be generated by the star during the early part of its Main Sequence life-time, but this makes only a small contribution to the heating of the cloud.     

Роль диффузного вешества В коронах галктик

I I I : M_gas<0.04 M_cor. - . , .     

High-dispersion spectroscopy of the symbiotic star AX Per: A binary model with rapid mass accretion

Results of high-dispersion spectroscopy (10 Å mm^–1) of the symbiotic star AX Per carried out in the years from 1979 to 1987 are reported. The emission line [FeVII] 6086 consists of a narrow and a broad component; the radial velocity of the narrow one varies according to the photometric period 681.6 days. This variation (K=30.6±1.5 km s^–1) seems to be due to the orbital motion of the hot star. The radial velocity of absorption lines varies with an inverse phase dependence and a much smaller amplitude (K=5.6±2 km s^–1), which may reflect the orbital motion of the red giant. The variation of the radial velocity of the emission lines of FeII, ect. (K=6.7±1.5 km s^–1) might be due to the rotation of the red giant. The profile of H emission line suddenly changed around the phase of the photometric minimum, which could be explained as a result of an eclipse of the emitting region by the red giant. On the other hand, some problems remain open in the behaviour of the radial velocities of H and HeI 5876.The observed results support a binary model of AX Per consisting of a rather massive (3M ) M-type giant and a Main-Sequence star (0.6M ). AX Per seems to be in an early stage of the Case C mass transfer, and the estimated very high mass accretion rate (10^–4 M yr^–1) is consistent with the theoretical models. The narrow component of the emission line of [FeVII] 6086 might be emitted in radiatively driven polar jets on the hot star of which luminosity is close to the Eddington limit.A new identification as ZrII at 6106.47 Å is proposed for the emission line at 6106 Å.     

The solar causes of geomagnetic disturbances

Geomagnetic disturbances have been identified with respect to their sources for 1977–1983. The disturbance level was found using the daily planetary index A _p. High-amplitude ( 50), mean-amplitude (24) and low-amplitude ( 12) disturbances are caused by solar flares of importance 1, coronal holes, and filament cavities, respectively. The ranges of probable amplitudes of disturbances of different nature and their relative number are found from Poisson random distributions of amplitudes.     

SERTS-97 Measurements Of Relative Wavelength Shifts In Coronal Emission Lines Across A Solar Active Region

We used slit spectra from the 18 November 1997 flight of Goddard Space Flight Center's Solar EUV Rocket Telescope and Spectrograph (SERTS-97) to measure relative wavelength shifts of coronal emission lines as a function of position across NOAA active region 8108. The shifts are measured relative to reference wavelengths derived from spectra of the region's nearby quiet surroundings (not necessarily at rest) because laboratory rest wavelengths for the coronal EUV lines have not been measured to sufficient accuracy for this work. An additional benefit to this approach is that any systematic uncertainties in the wavelength measurements are eliminated from the relative shifts by subtraction. We find statistically significant wavelength shifts between the spatially resolved active region slit spectra and the reference spectrum. For He ii 303.78 Å the maximum measured relative red shift corresponds to a Doppler velocity +13 km s^–1, and the maximum relative blue shift corresponds to a Doppler velocity –3 km s^–1. For Si x 347.40 Å, Si xi 303.32 Å, Fe xiv 334.17 Å, and Fe xvi 335.40 Å the corresponding maximum relative Doppler velocities are +19 and –14, +23 and –7, +10 and –10, and +13 and –5 km s^–1, respectively. The active region appears to be divided into two different flow areas; hot coronal lines are predominantly red-shifted in the northern half and either blue-shifted or nearly un-shifted in the southern half. This may be evidence that material flows up from the southern part of the region, and down into the northern part. Qualitatively similar relative wavelength shifts and flow patterns are obtained with SOHO/CDS spectra.     

Rapid simultaneous observations of Feii and Balmer emission lines of X Oph

Sudden brightening of FeII and Balmer (H and H) lines of X Oph was observed on 28–29 May, 1986. Equivalent width of FeII and H lines increased by a factor of two and that for H line by a factor of four, during the brigtening phase of the star. This brightening phenomena has been explained in the framework of Coronal Radiative Instability.     

The solar O i λ 7773 triplet

The reduction of observations of the O i 7773 triplet obtained at high spatial resolution (0.5) at two disk positions is described. Two sets of triplet profile data are presented at each disk position. One set represents data taken from 0.5 regions centered on the brightest granules, while the other set represents data taken from 0.5 regions centered on the cooler infalling intergranular material.     

Coronal environment of quiescent prominences

With thespectro-coronagraph and themultichannel subtractive double pass spectrograph (MSDP) at the Pic du Midi Observatory two quiescent prominences were observed simultaneously. From the spectro-coronagraph observations 2D maps of Hei 10830 , Fexiii 10798 and 10747 line intensities were obtained. In addition, we obtained 2D maps of the ratioR of the two iron lines. This ratio is used as a diagnostic for determining the density of the hot coronal plasma surrounding prominences. We found that the electron density is higher at the location of the prominences than in the corona, whereas small regions (40) of lower electron density are unevenly distributed around the prominences indicating that the surrounding corona is highly inhomogeneous. The density of the cavity is reduced by a factor 1.5 compared to the density of the prominence environment (5 × 10⁸ cm^–3). We discuss the existence of cavities around these prominences according to the orientation of their axes relative to the line of sight and according to the velocity field inside the prominences. Constraints on models for prominence formation are derived.     

Structure and Evolution of the Transition Region Network Observed in he ii λ304

The length scale and life time of the transition region network cells were studied using Heii 304 filtergrams. The temporal structure function was calculated from spatially aligned Heii 304 images. The estimated life time of the network cell was about 27 hr. We compared this life time with the life time of photospheric magnetic network and of the extrapolated magnetic network. The spatial structure function was calculated from the Heii 304 filtergrams. The calculated spatial structure function saturates at 25000 km. The transition region network elements are bigger in size than the photospheric magnetic network element. The magnetic network element equals the size of the Heii 304 network element when the photospheric magnetic field is extrapolated to a height of 3000 km above the photosphere where the magnetic fluxes are deployed. The derived value of the diffusion speed of the network elements was 0.098 km s^–1.     

Characteristics of enhanced and low-amplitude cosmic-ray diurnal variation

The occurrence of a large number of high- and low-amplitude cosmic-ray diurnal wave trains during the two solar cycles (20 and 21) over the years 1965–1990 has been examined as a function of solar activity. The high-amplitude days with the time of maximum in the 18:00 hr corotation direction do not indicate any significant correlation with solar activity. But, the low-amplitude days are inversely correlated with solar activity and the time of maximum shifts to earlier hours ( 15:00 hr direction). The slope of the power-specrum density roughly characterized by power spectral index n in the high-frequency range 3.5 x 10^–5 Hz to 8.3 x 10^–4 Hz (time scales of 20 min to 8 hr) is different for the two classes of events. A suggestion is made that the enhanced and low-amplitude cosmic-ray diurnal variations are produced by different types of interplanetary magnetic field distributions.     

CH3CN J = 8 − 7 and CS J = 3 − 2 emission in orion KL

The data of the line series CH₃CN 8(K) – 7(K) K = 0 – 7 and line CS J = 3 – 2 were taken simultaneously. At beam size of 16 the emissions of CH₃CN and CS have a common center position located near IRc2 with deviations -8 and 5. The observed data show that in Orion KL core the integrated intensities of the two species have double peaks separated by a space of 14. The 2-dimension Gaussian fitting plots (FWHM) are ellipses ofD _maj = 26 andD _min = 22 for CH₃CN 8(K) – 7(K) K = 3 – 6 andD_maj = 39 D_min = 31 for CS J = 3 – 2 at a distance about 450 pc. Towards the multiple line emission region of CH₃CN 8(K) – 7(K) K = 3 – 6, using a simplified very large velocity gradient model to solve the statistical equilibrium and radiative transfer equations, we find to fit the observed results, the optimum physical parameters and kinetic temperatureT _k 120 K, densityn(H _a) 1.2 × 10⁵ cm^–3, velocity gradientV _gr 92 km s^–1 pc^–1 and the local abundance of CH₃CNF _ab 3 × 10^–8. However towards the region of single line emission of CS J = 3 – 2 we have to use LTE and the optical thin approximation on the assumption ofT _k= 120 K to obtain the lower limits of column density and then, an averaged abundance of CS of 6 × 10^–8.     

Short-Term Periodicities in the Time Series of Solar Radio Emissions at Different Solar Altitudes

A spectral analysis of the time series of daily values of 12 parameters, namely, ten solar radio emissions in the range 275–1755 MHz, 2800 MHz solar radio flux, and sunspot numbers for six continuous intervals of 132 values each during June 1997–July 1999 showed considerable differences from one interval to the next, indicating a nonstationary nature. A 27-day periodicity was noticed in Interval 2 (26.8 days), 3 (27.0 days), 5 (25.5 days), 6 (27.0 days). Other periodicities were near 11.4, 12.3, 13.3, 14.5, 15.5, 16.5, 35, 40, 50–70 days. Periodicities were very similar in a large vertical span of the coronal region corresponding to 670–1755 MHz. Above this region, the homogeneity disappeared. Below this region, there were complications and distortions due to localized solar surface phenomena.     

Caii K line asymmetries in two well-observed solar flares of October 18, 1990

Two-dimensional evolutions of two flares of October 18, 1990 have been well observed in the Caii K line with a CCD camera at Norikura station of National Astronomical Observatory in Japan. There are two common characteristics for the flares: 3 - 5 min before the impulsive phase, the heating already begins at the footpoints of the flares, but no asymmetry in line emission has been detected. After the onset of the impulsive phase, Caii K line emission at the footpoints shows strong red asymmetry, with the maximum asymmetry occurring at the same time as the peak of the radio bursts. The maximum downward velocity is about 30 50 km s^–1. For flare 1, blue and red asymmetries were observed in two sides of the footpoint area. They developed and attained a maximum nearly at the same time and the inferred Doppler velocities are comparable (30 40 km s^–1). This implies that two mass jets started from a small region and ejected along a loop but in opposite directions with roughly equivalent momentum. A possible mechanism has been discussed.