On the intrinsic colours of cepheids

Some consequences from new data on the photometrically obtained intrinsic colours are considered. It is shown that the small amplitude and almost sinusoidal light-curve cepheids (Cs-subtype) increase their pulsation amplitude towards the red instability strip edge, differing from the other galactic cepheids. This feature is discussed in connection with Efremov's (1968) hypothesis that the Cs-cepheids first cross the instability strip from left to right. The galactic cepheid period-colour relation obtained by Dean, Warren and Cousins (1978) satisfies rather well the LMC cepheid observations and, consequently, considerations for the period-luminosity relations are made. The residuals from the PL relation proposed by us (AA line instead of Gascoigne's BB line in Figure 3) correlate with the colour residuals from the DWC period-colour relation V/(B–V) being equal to 2.7 (Figure 4). The luminosity effect as a possible cause of the descrepancy between the spectroscopic cepheid colours and the photometric colours is briefly discussed.     

Observational Study of the Three-Dimensional Magnetic Field Structure and Mass Motion in Active Regions

Spectro-polarimetric observations of active regions were carried out in the spectral lines of Sii 10827.1 Å and Hei 10830 Å to study the three-dimensional magnetic field structure and associated plasma flow properties. Comparison of Sii and Hei magnetograms with the potential field model shows that a large fraction of the magnetic field is consistent with the potential field structure, by assuming that the height difference between the origin of the two lines is about 1200 km. The slope of the scatter plot between Sii and Hei magnetograms is 0.5, 0.76 in an emerging flux and a larger active region, respectively. These values are lower than the scatter plot slopes obtained from Kitt Peak photospheric and chromospheric magnetograms, in which case the corresponding values are 0.83 and 0.9, respectively. Considering the height difference between these two sets of chromospheric magnetograms, this implies that the magnetic field spreads out faster near the transition region heights. Dopplergrams obtained by determining the centroid of the asymmetric line profiles show that, in case of emerging flux region, the chromospheric upflow regions are located in the magnetic neutral line areas.     

Semimajor axes of the orbits and ejection velocities of Perseid meteoroids

Photographic orbits of meteors are combined with modeling of the ejection of Perseids during the perihelion passage of comet Swift-Tuttle in 1862 to analyze the most likely ejection velocities of particles from the comet nucleus. Given the scatter of the semimajor axes of observed Perseids with masses greater than 10^–4 g, the most likely interval of ejection velocities spans 0 to 300 m/s for particles ejected in the plane of the comet orbit in the retrograde direction and in the direction of the comets anomalous tail.__________Translated from Astronomicheskii Vestnik, Vol. 39, No. 2, 2005, pp. 184–190.Original Russian Text Copyright © 2005 by Ishmukhametova, Kondrateva.     

Solar line blocking for disk-center and disk-averaged radiation from 3300 to 6860 Å

The line blocking is tabulated for 10 Å ( < 6300 Å) or 20 Å ( > 6300 Å) wide intervals. It follows from the spectral averages and the local continuum derived by Neckel and Labs from high-resolution Fourier transform spectra, which had been obtained by J. Brault at Kitt Peak. The internal accuracy (the scatter) is in the order of 0.1%. Significant systematic errors arising from local distortions of the adopted continuum level can be excluded. Larger errors are to be expected only near the Balmer limit, where the localization of the continuum is very ambiguous.     

Line formation in spherical media with partial frequency redistribution

Lines formed in a differentially expanding atmosphere have been calculated by using the angle averaged redistribution functionR _I (Hummer, 1962). We have compared these lines, in a few cases, with those formed by complete redistribution in the observer's frame of reference. We have considered an atmosphere whose ratios of inner to outer radii are 2 and 10, and it is assumed that the gas in the atmosphere is expanding uniformly with a maximum velocity of 3 mean thermal units. We have presented lines formed in spherical symmetry and those obtained by integrating over the apparent disc. Three types of physical situations are considered with a 2-level atom with non-LTE approximation (1) =10^–4, =10^–4, (2) =10^–4, =0 and (3) ==0, where is the probability per scatter that a photon is destroyed by collisional de-excitation, and is the ratioKc/Kl of absorption in the continuum per unit frequency interval to that in the line centre.It is found that there are noticeable differences between the profiles formed by partial redistribution (PRD) and complete redistribution (CRD). The profiles integrated over the stellar disc from the first type of the media are box type with flat top, which are similar to those observed in WC stars and those from the second type of media show emission peaks on red and blue sides with emission on the red side larger than the one on the blue side. The profiles from the third type of media with pure scattering, show some emission on the red side and deep absorption on the blue side. Large geometrical extensions of the atmosphere and higher gas velocities seem to enhance these two effects.     

Population I pulsating stars

Luminosities of Population I pulsating stars (Delta Scuti variables and classical cepheids) are investigated. From data for 80 Delta Scuti stars, semi-empirical period-luminosity-colour (P-L-C) relations and period-luminosity (P-L) relations are obtained for the four lowest modes of radial pulsations. The improvement of the accuracy of the stellar luminosity is determined when a P-L-C relation is used instead of the corresponding P-L relation. From data for 155 classical cepheids, empirical P-L relations are derived for short-period stars (logP1.1), long-period stars (logP>1.1), and s-cepheids. The comparison of the P-L relations for the two types of variable stars shows good agreement, but between them there is a gap with a dim nature.     

A novel look at the pulsar force-free magnetosphere

The stationary axisymmetric force-free magnetosphere of a pulsar is considered. We present an exact dipolar solution of the pulsar equation, construct the magnetospheric model on its basis and examine its observational support. The new model has toroidal rather than common cylindrical geometry, in line with that of the plasma outflow observed directly as the pulsar wind nebula at much larger spatial scale. In its new configuration, the axisymmetric magnetosphere consumes the neutron star rotational energy much more efficiently, implying re-estimation of the stellar magnetic field, \(B_{\mathrm{new}}^{0}=3.3\times10^{-4}B/P\), where \(P\) is the pulsar period. Then the 7-order scatter of the magnetic field derived from the rotational characteristics of the pulsars observed appears consistent with the \(\cot\chi\)-law, where \(\chi\) is a random quantity uniformly distributed in the interval \([0,\pi/2]\). Our result is suggestive of a unique actual magnetic field strength of the neutron stars along with a random angle between the magnetic and rotational axes and gives insight into the neutron star unification on the geometrical basis.     

Contribution à l'étude du spectre solaire de l'hydrogène

R'esumé Les raies et le rayonnement continu de l'hydrogène sont calculés, dans le cas solaire, en tenant compte des écarts à l'équilibre thermodynamique local.La comparaison des résultats, interprétés par la théorie de la formation des raies de Jefferies et Thomas, avec les observations donne des renseignements sur le modèle de l'atmosphère solaire.L'énergie émise par Ly dépend très fortement de la densité électronique. Seuls les modèles á palier de température dans la haute chromosphère donnent un profil de Ly à dépression centrale.La raie H, qui est contrôlée essentiellement pour les photoionisations dépend surtout du modèle de photosphère; toutefois les modèles de chromosphère à fortes température et densité électroniques ont une influence sur le profil de H.

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Summary The continuous and line spectrum of hydrogen emitted by the sun is computed, taking deviations of local thermal equilibrium into account. The theory of line formation as given by Jefferies and Thomas has been applied to the computations; a comparison of these results with the observations gives information on the solar atmospheric model.The energy emitted by Lyman has been computed for a two- and three-level atom: it appears to increase slightly with the number of levels. The energies, computed with the HAO model and the model of Coates, are larger than the observed values. The core of the Lyman profile is determined by collisions; the peaks are formed at a height where the electron temperature is about 20 000 °K; the energy depends very strongly on the electron density of the model.It is shown that the profile of Lyman only shows a central self reversal if the model of the high chromosphere has a temperature plateau.An interpretation of the observed distance of the peaks of the Lyman and Lyman line profiles is possible; it can also be shown why the distance of the Lyman peaks to the line centre is always of the order of 0.2 Å.The residual central intensity of H increases slightly with the number of atomic levels; the value computed with a five level atom, with the HAO chromospheric model, and the Utrecht photospheric model (1964) does not differ very much from the observed value: it is slightly smaller than the observed value and the computed profile is narrower than the observed profile.In the case of the HAO model the source function of H is dominated by the photo-ionization terms; nevertheless, the collision terms are not much smaller than the photo-ionization terms; the residual central intensity of H computed with a chromospheric model similar to the interspicular model of Athay and Thomas but assuming a higher electron temperature and density is larger than the observed value.

     

A minimization procedure applied to population synthesis in galaxy nuclei using a star cluster library: M31, M32

We apply the MINOS optimization system to the population synthesis of galactic nuclei, using a grid of star cluster equivalent widths as a function of age and metallicity. For some classes of red galaxy nuclei, this observational approach to population synthesis produces results similar to those predicted by the theoretical approach of Arimoto and Yoshii (1987). The synthesis results for blue nuclei tend to scatter more in the age vs metallicity plane, probably because we use only visible and near-infrared spectral data. Additional information at shorter wavelengths will possibly produce better-focused solutions in the plane. However, strong bursts of star formation can easily be identified. We provide in this contribution population synthesis for the central regions of M31 and M32. The bulge and the semi-stellar nucleus of M31 are dominated by the old metal-rich population: the semi-stellar nucleus has reached up [Z/Z ]0.6 and the bulge [Z/Z ]0.3. In the central region of M32, a metallicity up to the solar value has been reached and the synthesis indicates that the intermediate age component is not dominant.     

VLA and Trieste observations of type I storms,type IV and pulsations

A prime objective of this experiment was to determine whether type I or IV sources at 333 MHz contain features of small (arc sec) scale. With the VLA, our resolution was better than 4. However, we never observed any structure of size smaller than about 30, with the typical source sizes being between about 40 and 90.Many observations were simultaneous with the Trieste Astronomical Observatory records at 327 MHz. The observations were made on two days in November 1988. On 8 November the observations were of a type I storm about two hours after a major flare. On 14 November they were mostly of the main phase of a type IV event, including pulsations of a kind rarely seen, strongly circularly polarized, and having a well-defined period of about 12 s. The size of the pulsating source was about 40 by 60, and the brightness temperature was about 10⁹ K. We compare these pulsations with those observed earlier.     

Wavelength Shifts in the Solar Photospheric Spectrum

Wavelength shifts converted to velocities between solar lines observed at disc center and laboratory wavelengths of Fei, Feii, Tii, Nii, and Fei lines in the near infrared are plotted as a function of the logarithm of their solar equivalent width in milliångstroms. The need for wavelengths based on the wavelength standards is stressed. A comparison of photographic Fei solar wavelength is shown to agree, on the average, with Fourier Transform Spectrometer solar wavelengths within less than 0.5 milliångstroms. Using Balthasar's limb effect tables we convert the disc center velocities to limb velocities and find, though the scatter is large, that there is little evidence for a super-gravitational red shift.     

ON THE WAVELENGTH DEPENDENCY (AND ITS VARIATIONS) OF THE RATIO DISK-AVERAGED TO DISK-CENTER INTENSITY,F/I 0 (λλ0.385 TO 10 µm)

In the whole wavelength range, 0.385 to 10.0 µm, the ratios F/I ₀ of disk-averaged intensity F to disk-center intensity I ₀ follow very closely the simple relation F/I ₀ = 0.975 - 0.091^-1 - 0.00024^-5. The residuals exhibit a well-defined structure with minima and maxima in the order of± 5 × 10^-3 , which reflect the wavelength-dependency of the source function B(,) and the continuous absorption coefficient . The fact that the scatter of the coefficients of the relevant limb-darkening functions is usually much larger than the scatter of the F/I ₀ ratios seems to confirm the occurrence of limb-darkening variations, which however leave the F/I ₀ ratios almost unchanged. Nevertheless, minor differences between different observations, which concern the details in the wavelength-dependency of the F/I ₀ residuals, seem to indicate that even the F/I ₀ ratios undergo minor variations.     

Recurrent Population II X-ray transients — similarities to SU UMa cataclysmic variables

Quasi-periodic outburst activity is not uncommon among Population II X-ray binaries. This paper reports observations of such activity in several sources, made by the Vela 5B X-ray monitor. Typical periods are 1/2–2 years with an r.m.s. scatter in interval time of 10%. This activity is reminiscent of the superoutburst cycles of SU UMa CV's with respect to mean recurrence times, the variation of the recurrence times about their mean, and the total mass transferred during outburst. However, the outbursts in the X-ray sources have a substantially longer duration, 50–100 days instead of 10 days. I suggest that SU UMa and X-ray transient outbursts may be caused by similar mass-transfer instabilities.     

Large-scale convection cells in the solar atmosphere

The directions and velocities of meridional plasma motions are investigated using Doppler shifts of the magnetically non-split line Fe 557.6 nm. Possible drifts of the spectrograph were controlled by measuring nearly iodine lines from a laboratory source. The scattered light was kept low by using the Capri Coudé refractor mainly around local noon.There is a general scatter of about ±20 to ±50 m s^–1 in the yearly mean results. The scatter is up to –200 m s^–1 for the year 1985. The results are compared with published data. Although some systematic meridional plasma motions could be detected from the average of the meridional flows within the whole observing period 1982 until 1986, the final analysis suggests, that all meridional motions averaged over half a solar cycle are below ±10 m s^–1.Paper presented at the 11th European Regional Astronomical Meetings of the IAU on New Windows to the Universe, held 3–8 July, 1989, Tenerife, Canary Islands, Spain.This paper has also been presented at the IAU Colloquium No. 121Inside the Sun andAstronomy and Astrophysics 229, 224 (1990).     

Ionized helium in prominences and in the chromosphere

Quiescent prominences It is found that Heii 4686 is emitted in the same cold region of 10000 K as hydrogen, metal and neutral helium emission lines. This conclusion is based on the finding that the observed width of 4686 is the same as the calculated width of 4686. The calculated width is derived from the observed widths of hydrogen and metallic lines. The large intensity of Heii 4686 in 10000 K can be explained by the ionization of Heii due to the UV radiation below 228 Å that comes from the corona and the transition region.Loop prominences The very broad width (30 to 50 km s^–1) of 4686 for two post-flare loop prominences shows that the Heii line is emitted in hot regions different from regions of hydrogen and metal emission. From the widths of the Balmer lines and many metallic lines the kinetic temperature for one loop is found to be 16000 K in one part and 7600 K in another part. The electron densities are 10^12.0 cm^–3 and less than 10^11.0 cm^–3 respectively.Chromosphere The intensity of 4686 in the chromosphere can be interpreted in terms of a temperature of 10000 K with the ionization due to UV radiation. But, since observations of the width of 4686 are not available, a definitive conclusion for the chromosphere cannot be reached.     

Analysis of Perseid activity from short time visual count intervals

During periods of high meteor rates (> 100...400 per hour), visual observers can gather substantial samples in intervals of less than 20 minutes or so. This seems to allow to find real particle flux variations occurring in the meteoroid stream. Investigations of the 1993 and 1994 Perseid peak counts suggest that interval lenghts of 10 minutes are recommended for close to the peak periods. Perception effects increase the scatter of ZHRs obtained from < 10 minute counts. Thus flux variations occurring at scales of 50,000 km may be detected by visual 10 minute meteor counts.     

Simultaneous five color (UBVRI) polarimetry of VV Puppis

Observations of linear and circular polarization in five colour bands during a highly active state of VV Puppis in January-86 are reported. A strong linear polarization pulse with the maximum in the blue, P_B22 %, is observed at the end of the bright phase when the active pole is at the limb and a weaker secondary pulse, P_B7 %, is seen in the beginning of the bright phase, when the active pole reappears. Strong positive circular polarization is also observed in the blue and the ultraviolet, P_UP_B18 %, P_V10 % during the bright phase. The circular polarization reverses the sign in the B and V bands during the faint phase and a negative polarization hump is seen when the active pole crosses the limb. The circular polarization in the V band reaches the value P_V–10% at the hump, after which it remains near P_V–5% during the faint phase. This is probably due to radiation coming from the second, less active pole and accretion thus takes place onto both poles. The wavelength dependences of the positive and negative parts of the circular polarization curve are different and no polarization reversal is seen in the U band. The position angle of the linear polarization is well determined during a large portion of the cycle, especially in the V band, thanks to the activity from both poles. A best fit to the position angle curve, taking into account also the duration of the positive circular polarization phase interval =0.40 (in the V band), yields the values of orbital inclination i=78°±2° and the colatitude of the active magnetic pole =146°±2°. The relatively good fit to the position angle data indicates that the simple dipole model is nearly correct in the case of VV Puppis. Some wavelength dependence is, however, seen in the position angle curves, especially in the I band where the slope / at the main pulse is considerably smaller than in the other bands. The shape of the position angle curves changes also in the blue and the ultraviolet around the middle of the bright phase. This is probably due to optical thickness effects as the side of the accretion column which is toward the observer changes near this phase.Based on observations made at the European Southern Observatory, La Silla, Chile.Paper presented at the IAU Colloquium No. 93 on Cataclysmic Variables. Recent Multi-Frequency Observations and Theoretical Developments, held at Dr. Remeis-Sternwarte Bamberg, F.R.G., 16–19 June, 1986.On leave from University of Helsinki Observatory     

Profiles of the diffuse interstellar lines

The possibility that the diffuse interstellar lines and bands are but structure in the continuous extinction by dust grains is considered in detail. The lines are assumed to arise from impurities in the grains. Profiles of the strongest diffuse features are computed for a wide variety of host grain types and sizes, including size distributions of particles, and spheroidal and coated spherical grains.New observational profiles, of which 4430, 4765 and 4885 have been derived from automated spectrophotometry of electronographic spectra, are also presented; and comparison is made between theory and observation. The most useful tool of comparison is found to be the ratio of maximum apparent emission in the violet wing to maximum apparent absorption, . This ratio increases with size of the host particle, and with both real and imaginary parts of its refractive index; but is independent of the observed variations in strength and width of a particular line. The ratio is larger for 4430 than for the other lines.Taking into account current opinion as to the composition of interstellar dust, the observations may be explained by the presence of impurities in either silicate grains with effective radii of about 120 nm, or possibly graphite grains about 50 nm in radius coated with a thin dielectric condensate. None of the observed profiles are at variance with the hypothesis that all the lines arise from the same impurities in the same type of grains. The concentration of such impurities needed to reproduce the observed line profiles and strengths is only about one in 10⁴ atoms in the grain.     

Search for Spatial Variability in the Solar Acoustic Spectrum

Motivated by the various examples of spatial variability in the power of the acoustic spectrum, we attempted to look for spatial variability in the peak frequency of the spectrum. However, the determination of this peak frequency on a spatial scale of a single pixel (8 arc sec for the GONG data) is limited by the stochastic variations in the power spectrum presumably caused by the stochastic nature of the excitation process. Averaging over a large number of spectra (100 spectra from a 10 × 10 pixel area) produced stabler spectra. The peak frequencies of 130 such locations were found to be distributed with a FWHM of about 130 Hz. A map of the spatial variation of this peak frequency did not show any strong feature with statistically significant deviation from the mean of the distribution. Likewise, the scatter in the peak frequencies masked the detection of magnetic-field-induced changes in the peak frequency. On a much larger scale, the N latitudes showed a slightly lower value of the peak frequency as compared to the S latitudes, although the difference (25 Hz) is barely larger than the r.m.s. spread (20 Hz).     

Back-scattering of sunlight by ice grains in the mesosphere

Small dielectric ice particles of radius ~0.25 m, which are known to be present in the mesophere, scatter a fraction of incident sunlight in backward directions that do not reach the Earth. This back-scattered fraction is rigorously calculated using Mie theory for a uniform distribution of particles distributed over a sunlit hemisphere. Such calculations provide necessary information for estimating equilibrium surface temperatures of the Earth under different conditions.