Stationary equations of hydrogen excitation and ionization in prominences

The statistical equilibrium equations for the continuum and first 10 levels of a hydrogen atom show that the radiation of a bright prominence (the brightness of the H line has attained 56 mÅ of the disc centre spectrum) is completely due to scattering of the Sun radiation. The basic unknowns are separated with certainty: electron concentration (n _e = 3.0 × 10¹⁰ cm^–3), effective thickness (l = 4.2 × 10⁸ cm) and electron temperature (T _e = 5000 K).Radiation of a very bright prominence (A (H) = 213 mÅ; T _e = 7300 K; n _e = 5.0 × 10¹¹ cm^–3; l = 1.3 × 10⁷ cm) is on account of electron impacts (40%) and the Sun radiation scattering (60%).The parameters are shown to depend greatly on the prominence optical thickness in the lines of the first subordinate series of a hydrogen atom. In the course of determination all the parameters and 100 interconnected integral equations of the radiation diffusion have been thickness-averaged; the population of levels has been calculated by observations using the self-absorption factors.     

Time variability and structure of quiet Sun sources at 6 cm wavelength

Using the Westerbork Synthesis Radio Telescope (WSRT) we produced a synthesized map of a quiet Sun region on June 15, 1976, and studied the structure and time variability of the quiet emitting regions at 6 cm wavelength with a spatial resolution of 6 arc sec. Comparison of the 12hr synthesis map with Ca⁺ K filtergram shows that bright and dark features on the 6 cm quiet Sun synthesized map correspond to the chromospheric networks and cells observed in Ca⁺ K. All 6 cm bright features lie over bright Ca⁺ K network elements. The reverse correlation is not true, that is, not all bright Ca⁺ K network features have their 6 cm counterparts. Comparison with the photospheric magnetogram shows that about 72% of the photospheric magnetic field enhancements (¦B¦ 5 G) are coincident with 6 cm emissive regions. Only one 6 cm feature could be positively identified with a bipolar magnetic structure. This implies that no more than 20–25% of the 6 cm emitting features could be associated with X-ray bright points. Intercomparison of our 12hr two-dimensional synthesis map, a 4hr two-dimensional synthesis map (around meridian) and the one-dimensional fan beam scans of the quiet Sun region at 6 cm, along with the Ca⁺ K filtergram and photospheric magnetogram shows that: (1) All of the 15 time-varying elements at 6 cm were located on Ca⁺ K networks; (2) about 40% of the 15 time varying elements at 6 cm are coincident with enhancements of the photospheric magnetogram; (3) individual time-varying sources have minimum source size (FWHM) of 15 arc sec and maximum brightness temperature of 10⁵ K; (4) the life time of the time varying sources varies from a few minutes to several tens of minutes; (5) the intensity of the sources varies by factors of 2 to 7 over time periods of 1 min to tens of minutes; and (6) the sources tend to disappear for periods of up to tens of minutes and to reappear at the same locations.     

Chromospherically active binaries,kinematics and age connections

Space velocities of 146 chromospherically active binary stars have been calculated. Containing F to M spectral types on the Main Sequence together with G and K giants, this very heterogeneous sample has been divided into subsamples in order to segregate stars so that they have similar kinematics and ages. Dispersions of space velocity components and other kinematical quantities (velocity averages and |Z| distributions) of these groups imply their ages as: Both old and young Main-Sequence systems (<10×10⁹ yr and 1×10⁹ yr, respectively) exist in the sample. Systems containing subgiants (single- and double-lined) or double-lined giants as their companions have ages about 2–3×10⁹ yr. Single-lined giants appear to be older than intermediate disk population stars (>5×10⁹ yr). The possible existence of white dwarfs as invisible companions of some of the single lined giants is suggested in order to explain why these systems are older than double-lined giants.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.     

Excitation of Orbital Inclinations of Asteroids during Depletion of a Protoplanetary Disk: Dependence on the Disk Configuration

The pumping up of orbital inclinations of asteroids caused by sweeping secular resonances associated with depletion of a protoplanetary disk is discussed, focusing on the dependence on the disk inclinations and surface density distribution. The asteroids have large mean inclinations that cannot be explained by present planetary perturbations alone. It has been suggested that the sweeping secular resonances caused by disk depletion are responsible for these high inclinations. Nagasawa et al. (2000, Astron. J.119, 1480-1497) showed that the inclinations of asteroids are pumped up if the disk is depleted in an inside-out manner on a time scale longer than 3×10⁵ years. Their assumed disk midplane is not on the invariant plane. However, it should be affected by the inclination of the disk plane. Here we investigate the dependence on the disk inclinations. We assume a disk depletion model in which the disk inside the jovian orbit has been removed and the residual outer disk is uniformly depleted. We calculate the locations of the secular resonances and the excitation magnitude of the inclinations with analytical methods. We found that the inclinations are pumped up to the observational level for a depletion time scale longer than 10⁶ years in the case of the disk plane that coincides with the invariant plane. The required time scale is longest (3×10⁶ years) if the disk plane coincides with the jovian orbital plane. However, it is still within the observationally inferred depletion time scale. We also studied dependence on a disk surface density gradient and found that the results do not change significantly as long as the inner disk depletion is faster than the outer disk one.     

The Hα development of an intense limb flare and associated flaring arches

The H analysis of the development of the strong impulsive and faint gradual phase of the June 26, 1983 flare indicates the following: (1) The flare originated from two microprominences on the southeast border of NOAA 4227. Several similar events are summarized in Table II. (2) The main flare structure was a flare cone, which consisted of a bright surge-like stream, elevated above two flare ribbons (located in the cone's base). The flare cone had a height of about 40 × 10³ km and lasted 4 min in H. The upper part of the cone was terminated by a very fine loop, which was bent to the west, where later a chromospheric brightening occurred at the footpoint of a flaring arch. A 300 keV burst and radio spikes were observed during the maximum flare phase. (3) The flaring arch system, with its apex at a height of about 48 × 10³ km, formed the skeleton for the coronal helmet structure (Figure 7(c)). The velocity of the plasma moving along the flaring arch was between 3500 km s^–1} and 6900 km s^–1} during the first brightening (14:07 UT).     

Measurements of the helium 584 Å airglow during the Cassini flyby of Venus

The helium resonance line at 584 Å has been observed with the UltraViolet Imaging Spectrograph (UVIS) Extreme Ultraviolet channel during the flyby of Venus by Cassini at a period of high solar activity. The brightness was measured along the disk from the morning terminator up to the bright limb near local noon. The mean disk intensity was ∼320 R, reaching ∼700 R at the bright limb. These values are slightly higher than those determined from previous observations. The sensitivity of the 584 Å intensity to the helium abundance is analyzed using recent cross-sections and solar irradiance measurements at 584 Å. The intensity distribution along the UVIS footprint on the disk is best reproduced using the EUVAC solar flux model and the helium density distribution from the VTS3 empirical model. It corresponds to a helium density of 8×10⁶ cm⁻³ at the level of where the CO₂ is 2×10¹⁰ cm⁻³.     

Stellar pulsation: (I) analysis of stability of envelope

The life-time of the star on AGB is approximately 6 × 10⁴ yr. We divide it into front half and back half of AGB (including to optical Mira variable and OH/IR star) according to their evolution character. The observations show that the star has non-pulsation, but constant mass loss rate ( 5 × 10^–7 M yr^–1) on front half of AGB. Its circumstellar envelope is formed. When the star has pulsation on back half of AGB, its mass loss rate is relative with time, and increases gradually. In this time the star is on the stage of optical Mira variable. When the mass loss rate reaches the value of 3 × 10^–6 M yr^–1, the star evoluted from the stage of optical variable into the stage of radio bright OH/IR star. On the end of AGB the mass loss rate reaches 10^–4 M yr^–1. (Band and Habing 1983, Hermen and Habing 1985).     

Observations of prominences at 3.5 millimeter wavelength

At 3.5 mm wavelength absorption features are observed in correspondence with H dark filaments on the disk; beyond the limb the prominences correspond to emissive regions. The absorption features are larger (2–3 arc) than the corresponding H dark filaments; the emissive regions at the limb have similar angular sizes. The emissive regions at the limb have electron temperatures of 5500±500 K; the amount of absorption observed on the disk leads to mean electron densities of about 5 × 10¹⁰ per cm³.     

Lα, Lβ (of Hi), k and h (of Mgii), K and H (of Ca ii) observations in a quiescent prominence with the OSO-8 LPSP instrument

A sequence of images taken at different positions in the resonance lines of Ca ii, Mg ii, and H i was obtained over a quiescent prominence with the LPSP instrument on OSO-8. Ca ii K (and H) profiles are reconstructed at different locations in the prominence with a (10 × 5) arc sec² resolution. Significant variations of FWHM and line shifts are found: FWHM range from 0.14 Å to 0.5 Å; blue shifts reach about 14 km s^-1. The ratio of K to H absolute intensities shows a large spread around the average value of 1.2. The same ratio for the Mg ii lines in the whole prominence is higher (1.7), a fact already noticed at the edge of an active prominence (Vial et al., 1979). The ionization degree, as measured by the L/Ca K ratio, shows noticeable variations within the prominence. The L intensity is about 0.3 times the intensity measured in the quiet Sun, and the L/L ratio is less than one half the disk value. These results indicate important variations of the thermal conditions inside the prominence.DASOP, Observatoire de Paris, 92190 Meudon, France.     

Activity in the quiet Sun

Macrospicules have been observed in H and He i D₃, on the disk and above the limb. In 1975, a rate of 1400 (A day)^–1 is inferred, and the ratio of equatorial to polar rates 2. D₃ intensities are a few × 10^–3 of the disk center, and do not decrease in coronal holes. The ratio of H to D₃ intensities is 10. The integral number of macrospicules with D₃ intensity I ₀ is proportional to I ₀ ^–1.     

Studies of solar flares using optical,X-ray and radio data

I have studied a number of flares for which good X-ray and optical data were available. An average lag of 5.5 s between hard X-ray (HXR) start and H start, and HXR peak and Ha peak was found for 41 flares for which determination was possible. Allowing for time constants the time lag is zero. The peak H lasts until 5–6 keV soft X-ray (SXR) peak. The level of H intensity is determined by the SXR flux.Multiple spikes in HXR appear to correspond to different occurrences in the flare development. Flares with HXR always have a fast H rise. Several flares were observed in the 3835 band; such emission appears when the 5.1–6.6 keV flux exceeds 5 × 10⁴ ph cm^-2 s^-1 at the Earth. Smaller flares produce no 3835 emission; we conclude that coronal back conduction cannot produce the bright chromospheric network of that wavelength.The nearly simultaneous growth of H emission at distant points means an agent travelling faster than 5 × 10³ km s^-1 is responsible, presumably electrons.In all cases near the limb an elevated Ha source is seen with the same time duration as HXR flux; it is concluded that this H source is almost always an elevated cloud which is excited by the fast electrons. A rough calculation is given. Another calculation of H emission from compressed coronal material shows it to be inadequate.In several cases homologous flares occur within hours with the same X-ray properties.Radio models fit, more or less, with field strengths on the order of 100G. A number of flares are discussed in detail.     

An Erupting Active Region Filament: Three-Dimensional Trajectory and Hydrogen Column Density

From 15:33 through 16:02 UT on 13 June 1998, observations of an erupting filament as it crossed solar disk center were obtained with the NSO/KPVT and SOHO/CDS instruments as part of the SOHO Joint Observing Program 70. Context observations show that this event was the eruption of the north-east section of a small active region filament associated with NOAA 8237, that the photospheric magnetic field was changing in this active region between 12–14 June 1998, and that a coronal Moreton-wave disk event occurred, as well as a white-light CME off the south-west solar limb. The NSO/KPVT imaging spectroscopy data covered 512 × 512 arc sec of the disk center and were spectrally centered at the Hei 1083 nm line and captured ±1.0 nm of surrounding solar spectrum. The Hei absorption line is seen blue-shifted to velocities of between 200 and 300 km s^–1. The true solar trajectory of the eruption is obtained by using the projected solar coordinates and by integrating the Doppler velocity. The filament travels with a total velocity of about 300 km s^–1 along a path inclined roughly 49 deg to the solar surface and rises to a height of just over 1.5 solar radii before it becomes too diffuse to follow. The filament also shows internal motions with multiple Doppler components shifted by ±25 km s^–1. Finally, the KPVT data show no Stokes V profiles in the Doppler-shifted Hei 1083.03 nm absorption to a limit of roughly 3×10^–3 times the continuum intensity. The SOHO/CDS scanned the center of the KPVT FOV using seven EUV lines; Doppler-shifted filament emission is seen in lines from Hei 58.4 nm, Heii 30.4 nm, Oiv 55.5 nm, Ov 63.0 nm, Nevi 56.3 nm, and Mgx 61.0 nm representing temperatures from about 2×10⁴K through 1×10⁶K. Bound-free continuum absorption from Hi, without confusion from foreground emission and line emission, is seen as the filament obscures underlying chromospheric emission. A fit to the wavelength dependence of the absorption from five lines between 55.5 to 63.0 nm yields a column density _H I =4.8±2.5×10¹⁷ cm^–2. Spatial maps show that this filament absorption is more confined than the regions which show emission.     

Three-wave turbulence in pulsar winds

Three-wave interactions can modify the very bright radio waves propagating through dense, magnetized pulsar winds and prevent radiation from escaping, or severely alter the resulting spectrum. Absence of the signatures of such effects in the very bright radio emission from pulsars implies constraints on the structure of their winds, especially for a spherically-symmetric wind which should block the radiation unless that Lorentz factor of the pairs is 8 × 10².Presented at the Cosmic Winds and the Heliosphere conference in Tucson, Arizona, U.S.A., October 18–22, 1993.     

The accretion matter at the disk rim in Her X-1/HZ Her

The Crosa and Boynton (1980) empirical model for discrete mass transfer in Her X-1 is further developed. The photometric features of the light curve (peaks of an hour duration and 0.3–0.7 ^m amplitude, steps near orbital phase =0); and the linear polarization bursts are assumed to be due to the formation and eclipses of the plasma blobs produced by discrete transfer of matter from optical star surface and its interaction with the accretion disc rim. The long lifetime (20^h) of the cold (3×10⁴ K) blob extending up to 10¹¹ cm above the disc plane, as well as the deep X-ray flickerings (300 s) during the X-ray absorption dips are assumed to arise from a dispersal of accreting matter by the Rayleigh-Taylor instability in a blob moving through a hot corona of the disk atT _c =3×10⁶ K andn _c =3×10¹¹ cm^–3. Thermal equilibrium in the corona and in the blobs are supported by X-ray flux. Within the first few hours after its formation a blob disintegrates into drops withr=5×10⁹ cm,T=3×10⁴ K, andn=3×10¹³ cm^–3 which move then along Keplerian orbits. Frictional interactions of the drops with the corona destroy them on a 20^h time-scale. The proposed model makes it possible to interpret the diverse observational facts and to predict numerous observational displays in the optical, UV, and X-ray bands. The first results of our optical-spectrum observations of blobs are briefly described.     

Properties of the noise storm source at wavelength 1 m from observations of the solar eclipse on March 29, 2006

Observations of the solar eclipse on March 29, 2006, at the Laboratory of Radio Astronomy of the CrAO showed that the radio radius of the Sun at a wavelength of 1 m in the direction of the first contact was R _d = 1.12 R _⊙ during solar activity minimum between cycles 23 and 24. The brightness temperature of the undisturbed Sun was T _d = (0.6 ± 0.06) × 10⁶ K. There was a noise storm source above the sunspot group NOAA 0865 whose bright nucleus had a size of 1′.3 and a brightness temperature T _b = 16 × 10⁶ K. The noise storm bursts were emitted from the region of the bright nucleus above the group NOAA 0865 and were absent during its covering by the disk of the Moon. Thermal radiation from a coronal condensation with a brightness temperature of (1?2) × 10⁶ K extending out from the visible solar disk to 2′.7 was observed during the eclipse above the eastern limb sunspot group NOAA 0866. The bright nucleus in this limb source appeared 42 min after eclipse termination and persisted in the ensuing days. This may be indicative of the time of its emergence from behind the radio horizon formed by regular refraction of radio waves in the corona. The refractive displacement was measured by comparison with the eclipse observations at a shorter wavelength of 12 cm. Its value of 0′.96 is close to the calculated value of 0′.8.     

Resonant structure of the outer Solar System

Hierarchical stability of the outer Solar System is monitored through its 3-body subsystems by using numerically computed ephemerides for 5×10⁶ yr. It is found that the stability parameters of Sun-Jupiter-Saturn and Sun-Uranus-Neptune oscillate in anti-phase in 1.1×10⁶ yr. The mechanism responsible for this locking is a secular resonance between Uranus' perihelion and Jupiter's aphelion: the difference between the two librates within ±70° with the same period of 1.1×10⁶ yr.     

Rocket-coronagraph photometry of the 7 March, 1970 corona from 3 to 8.5 R s

A rocket-borne coronagraph utilizing external occulting disks was used to photograph the solar corona from 3 to 9 R _s at 1931 UT on 7 March, 1970. Comparison of the rocket and ground-based observations shows a one-to-one correspondence between major streamers from the inner to the outer corona. In particular streamers over the poles are clearly visible against the background corona from 3 to 8 R _s. These rocket data had a scattered light level of 1.2 × 10^–10 B _s. The derived quiet equatorial and polar K + F corona was within 10% of the absolute brightness of standard coronal models and displayed identical radial gradients to those models. The photometric profiles of the NE limb streamer were analyzed assuming a model in which the core density follows a Gaussian distribution in directions perpendicular to the radius vector. This streamer was assumed to be rooted on the visible disk at 55 to 60° from the plane-of-the-sky as based on K-coronameter and XUV data. An uncertainty of as much as a factor of three still remains for the value of the axis density owing to uncertainty in the line-of-sight dimension of the streamer.     

Coronal bright points at 6 cm wavelength

We report the results of the first observations of solar coronal bright points at 6 cm wavelength using the Very Large Array (VLA), with a spatial resolution of 1.2. The maximum brightness temperature of the sources observed is 3 × 10⁴ K with a mean value of 1 × 10⁴ K (above the quiet Sun value). The lifetime of most sources is between 5 and 20 min. The average diameter of the sources is about 5–15 arc. The sources are gaussian-like near the footpoint of miniature loops and they appear in groups. The observations indicate that significant fluctuations in the brightness temperature (sometimes quasi-periodic) and in the spatial extents of these sources can occur over periods of a few minutes.On leave from Beijing Observatory, Beijing, Peoples Republic of China.     

Kinematics and parameters of the gas in the vicinity of TW Hya

The following conclusions about the kinematics and parameters of the gas in the vicinity of TW Hya have been drawn from an analysis of optical and ultraviolet line profiles and intensities. The accreting matter rises in the magnetosphere to a distance z>R_* above the disk plane and falls to the star near its equator almost perpendicular to its plane. The matter outflows from a disk region with an outer radius of ≤0.5 AU. The [OI], [SII], and H2 lines originate in the disk atmosphere outside the outflow region, where the turbulent gas velocity is close to the local speed of sound. In the formation region of the forbidden lines, T?8500 K and N_e?5×10⁶ cm^?3, and the hydrogen is almost neutral: x_e<0.03. The absorption features observed in the blue wings of some of the ultraviolet lines originate in the part of the wind that moves almost perpendicular to the disk plane, i.e., in the jet of TW Hya. The V _z gas velocity component in the jet decreases with increasing distance from the jet axis from 200 to 30 km s^?1. The matter outflowing from the inner disk boundary, moves perpendicular to the disk plane in the formation region of blue absorption line components, at a distance of ～0.5 AU from the axis of symmetry of the disk. This region of the wind is collimated into the jet at a distance of <3 AU from the disk plane. The gas temperature in the formation region of absorption components is ?2×10⁴ K, and the gas density is <3×10⁶ cm^?3. This region of the jet is on the order of several AU away from the disk plane, while free recombination in the jet begins even farther from the disk. The mass-loss rate for TW Hya is \(\dot M_w < 7 \times 10^{ - 10} M_ \odot yr^{ - 1}\), which is a factor of 3lower than the mean accretion rate. The relative abundance of silicon and aluminum in the jet gas is at least an order of magnitude lower than its standard value.     

Observations of Ellerman bombs in H α

A developing active region near the center of the solar disk was observed for 80 min at the center and the wings of H. Ellerman bombs lying both below an Arch Filament System and near sunspots were studied at H - 1.0 Å and H - 0.75 Å. We determined their average contrast, lifetime, size and we studied their flux as a function of time. We found evidence that the size of Ellerman bombs increases with height. The time curves of flux provide evidence for both impulsive and gradual energy release. Under the AFS the Ellerman bombs form a cellular pattern with a characteristic size of 3.1 × 10³ km. Fifty percent of the bombs appear and disappear in pairs, possibly associated with bipolar emerging magnetic flux tubes.