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.     

The supergranule velocity field

A study of supergranule motions confirms horizontal velocities with peak values of typically 0.36 km s^–1 as observed in Fe i 8688 Å. These show no significant variation with height over the range of formation of C i 9111, Fe i 8688, and Mg i 8806, but there is a substantial reduction to about one-half of this at the level of Ca ii 8542.Near disk center, supergranule vertical velocities in Fe i 8688 have rms values ±0.01 km ^–1, after allowance for the residual effects of the line-of-sight component of the horizontal supergranule motions, the five-minute oscillations, granule motions, and detector drift. There is a marginally-significant association of magnetic elements, and hence of cell boundaries, with downward motions; but this requires further testing.Measurements of downward velocities 0.1 km^–1 in regions of strong magnetic field when using unpolarized light are attributed to the much higher downflow inside the elements themselves and have nothing to do with supergranule motions.Visiting Astronomer, Kitt Peak National Observatory.Operated by the Association of Universities for Research in Astronomy, Inc. under contract with the National Science Foundation.     

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.     

Circumstellar line of Cai λ6572 in Zeta Aurigae and 31 Cygni

In the 1974 eclipse of Zeta Aurigae the satellite line of the Cai intercombination line at 6572 Å appeared in all the phases observed, i.e., near the mid-eclipse and egress, although its equivalent width, 100–200 mÅ, and velocity deviation, –20 im s^–1, from the principal line were considerably less than those for the satellite line found just after the fourth contact of the 1971–72 eclipse. 31 Cygni also showed a similar satellite line to Zeta Aurigae's in the 1974 eclipse, not only in the 1972 eclipse but also outside the eclipse. These satellite lines should be due to the circumstellar gas expanding from the binary systems.     

O VI (λ = 1032 Å) profiles in and above an active region prominence,compared to quiet Sun center and limb profiles

O vi ( = 1032 Å) profiles have been measured in and above a filament at the limb, previously analyzed in H i, Mg ii, Ca ii resonance lines (Vial et al., 1979). They are compared to profiles measured at the quiet Sun center and at the quiet Sun limb.Absolute intensities are found to be about 1.55 times larger than above the quiet limb at the same height (3); at the top of the prominence (15 above the limb) one finds a maximum blue shift and a minimum line width. The inferred non-thermal velocity (29 km s^–1) is about the same as in cooler lines while the approaching line-of-sight velocity (8 km s^–1) is lower than in Ca ii lines.The O vi profile recorded 30 above the limb outside the filament is wider (FWHM = 0.33 Å). It can be interpreted as a coronal emission of O vi ions with a temperature of about 10⁶ K, and a non-thermal velocity (NTV) of 49 km s^–1. This NTV is twice the NTV of quiet Sun center O vi profiles. Lower NTV require higher temperatures and densities (as suggested by K-coronameter measurements). Computed emission measures for this high temperature regime agree with determinations from disk intensities of euv lines.     

9Be abundances in dwarfs of intermediate metal deficiency: Impications for galactic evolution

We present observations of the 3130 Å Beii resonance doublet in stars of intermediate metal deficiency, –0.6[Fe/H]–1.1 obtained with the Intermediate Dispersion Spectrograph and IPCS detector at the 2.5 m Isaac Newton Telescope on La Palma. The derived beryllium abundances range from 5.6×10^–12 (one hafl solar) to 2×10^–12 (one sixth solar). These values, interpolated between the sets of previous beryllium measurements at higher and at lower metallicities, serve to confirm the monotonic variation of the Be abundance with Fe during the evolution of the galactic disc. We find that there was no extreme burst of Be production in the halo. To circumvent the effects of depletion, a statistical set of data is needed, for which the upper envelope can be used to trace galactic Be evolution. We note that no observations with sufficient sensitivity to detect Be in the extremely metal deficient stars of the halo have been reported to date.     

Na i D brightening phenomena in quiescent prominences

The Na i D emission lines are found to brighten temporarily in restricted regions of quiescent prominences and we call the phenomena the Na i D brightenings. The comparison of observed intensities of Na i D lines with model calculations shows that the phenomena is attributed to a kind of local activation of quiescent prominences accompanied by mass motions. The Na i D lines are emitted from extraordinarily high pressure regions in which the pressure rises up to 0.37 2.7 dyn cm^–2 and the temperature seems to be in the range between 5800 7000 K. The line-of-sight velocity of the mass motions amounted to several tens of km s^–1 in some Na i D emitting prominences investigated. The life time of the phenomena is estimated to be about several tens of minutes.     

Physical parameters of reflection nebulae in the galaxy

We show how, given observed equivalent widths of Mgii and Mgi absorptions due to an interstellar cloud in which a late-B star is embedded, the basic physical parameters: kinetic temperature, mean density, electron density, and radius can be constrained. Hydrogen ionization by means of cosmic rays and the effect of the stellar radiation field on the magnesium ionization equilibrium are taken into account.The method is applied to the reflection nebula surrounding the star HD 26676. The resulting solutions for the radius and temperature of the nebulosity are comparable to the typical values derived for diffuse interstellar clouds from optical and 21-cm measurements, if a cosmic-ray ionization rate 10^–16s^–1 — in agreement with recent determinations — is assumed. The results are not strongly dependent on the gas pressureP forP varying in a range of values typical of interstellar clouds.     

Profiles of Hi (La), Mgii (h and k), Caii (H and K) Lines of an active filament at the limb,with the LPSP instrument aboard the OSO-8 Satellite

We scanned the H i L, Mg ii h and k, Ca ii K and H lines simultaneously with the LPSP instrument on OSO-8, to investigate the low and moderate temperature regions of an active region filament. The L line is not reversed except for the innermost position in the prominence. Intensity (k/h), (K/H) ratios are respectively 2 and 1.1, indicating that the Mg ii lines are optically thin, and that Ca ii K is saturated, although not clearly reversed. The results obtained during the second sequence of observations (K saturated before L for example) indicate that within the size of the slit (1 × 10) we are not observing the same emitting features in the different lines.We also observe an important line-of-sight velocity at the outer edge of the feature, increasing outwards from a few km s^–1 to 20 km s^–1 within 2. Less than half an hour later, this velocity is reduced to 15 km s^–1 while the intensities increase. Full width at half maximum intensities for this component indicate turbulence variations from 22 to 30 km s^–1. The observed high velocities at the top of the prominence can be compared with radial velocities that Mein (1977) observed in H at the edges of an active filament and interpreted as velocity loops slightly inclined on the axis of the filament.     

The Formation of MgI 4571 Å in the solar atmosphere

The two-dimensional equation of transfer is solved for the case of locally-controlled source function (LTE) and radiationally-controlled ionization. Horizontal fluctuations in electron temperature and macroscopic velocity fields are superposed on the basic one-dimensional model (cf. Altrock and Cannon, 1972). Output intensities are compared with observed rms intensity fluctuations and spatially-averaged intensities in Mg i 4571 Å. We find that at least one model (with a height-independent temperature fluctuation T/T=±0.02 in the range 0h450 km) can predict the magnitude of the intensity fluctuations in both the continuum and 4571 Å. The asymmetry of the line can be explained by adding a height-independent, temperature-correlated flow of amplitude 1 to 2 km s^–1. The relationship between these results and other multi-dimensional analyses is discussed.On leave from Department of Applied Mathematics, University of Sydney, Sydney, Australia.     

Circumstellar clouds derived from the 2800 Mgii data

As a result of the analysis of the observed interstellar 2800 Mgii absorption line data, an empirical relationship — a positive correlation — between the equivalent widthW(2800) and the effective temperature of the starT was discovered (Figure 1). However, in the case when this doublet is of stellar (photospheric) origin, only a negative correlation betweenW(2800) andT exists. Hence, the existence itself of such a positive correlation betweenW(2800) andT may be viewed as incomprehensible for the present influence of the star on the strength of the absorption line 2800 Mgii of nonstellar origin.On the other hand, we have evidence that the ionizing radiation of hot stars cannot provide for the observed very high degree of ionization of the interstellar magnesium. In particular, the observations give for interstellar magnesium the ratioN ⁺/N ₁ 1000, while in the case of ionization under the action of stellar radiation only we haveN ⁺/N ₁ 10.The assumption that circumstellar clouds surround hot stars can naturally explain these and other similar facts. A method for the determination of the general parameters-size, concentration, mass etc. — of the circumstellar clouds is developed. The main results of the application of this method to the relation of more than 20 hot stars are:(1) The circumstellar clouds surround almost (70%) all hot giants and subgiants. In the remaining (30%) cases, the absence of circumstellar envelopes requires additional evidence. (2) The linear sizes of circumstellar clouds vary within wide ranges — from 0.002 pc up to 1 pc. Most frequent are clouds with size of 0.1 pc. (3) The main concentration of hydrogen atoms (electrons) in circumstellar clouds is of the order of 100 cm^–3; the minimum value is 20–30 cm^–3, the maximum 10⁴ cm^–3. In one case (Deneb) the electron concentration rises up to 10⁵ cm^–3 for the size of the cloud 0.001 pc=3×10¹⁵ cm. (4) Stars of the same spectral and luminosity classes may possess circumstellar clouds characterized by quite different parameters. (5) Hydrogen in circumstellar clouds is completely ionized; for these clouds the optical depth _c 1; on the average,T _c 0.005. (6) The integrated brightness of circumstellar clouds is substantially fainter (by 8–10^m) than that of the central star. This is the reason why these clouds cannot be detected by ground-based observations. (7) The masses of individual circumstellar clouds vary from 1 down to 10^–4 . This gives for the mass ejection rate from 10^–10 to 10^–6 per year in case if these clouds are formed by the braking and accumulation of the ejected mass.The method of 2800 Mgii seems very convenient, fruitful and promising for the detection and study of circumstellar envelopes. Also, this method is very sensitive for a determination of the general parameters of such clouds, and concerns practically all their geometric, physical, kinematic and other properties.     

Brightness fluctuations in the K-line wings

A power-spectrum and cross-spectrum analysis has been made of measurements of temporal fluctuations of intensity observed in the K-line wing (2.07 Å from line center) and of simultaneous measurements of temporal fluctuations of Doppler displacement of the cores of 3931.122 Fe i and 3933 Ca ii (K₃). The measurements were made in a quiet region near the center of the Sun's disk. We find that the average power spectra of the intensity fluctuations have two significant peaks of about equal strength: one at 0.0033 Hz (300-s period); and one at about 0.001 Hz (1000-s period). The average rms value of these intensity fluctuations is 0.0435±0.0082. Maximum brightness comes before maximum violet displacement of the Fe i line. The mean of the best determined phases is 137° and of all the data 108°. At those places on the Sun where the 300-s oscillations can be identified in the k₃ core, the Doppler displacement of the Fe i line leads that of the K₃ core by a mean phase angle of 27°.     

The spectrum of the Ap star HR 465 near the rare-earth minimum phase

Four spectra of the long-period Ap star HR 465 were taken in 1969–70 near its rare-earth minimum phase and analysed. Equivalent widths or upper limits were obtained for lines of 23 elements. Lines of Cr were very strong. The relative abundances of the iron-peak elements were derived by a differential curve-of-growth analysis using Gem as the comparison star. The Sc/Fe, Ti/Fe, V/Fe, and Ni/Fe abundance ratios were found to be abnormally low, while the Cr/Fe ratio was abnormally high (0.08). The equivalent widths in 1969–70 were compared with those measured by Aller on a spectrum taken in 1960 near the rare-earth maximum phase. Ionized lines of Sc, Ti, V, Te, Nb, Mo, and the rare-earths all vary out of phase with the Cri and Crii lines. Although Mo and Cr occupy corresponding positions in the periodic table, the Mo/Cr abundance ratio was apparently at least 180 times higher in 1960 than in 1969–70.     

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.     

A spectroscopic study of solar spicules in Hα, Hβ and K

Simultaneous spectra of spicules in H, H and the Ca ii K have been analyzed. These spectra were obtained with the échelle spectrograph of Sacramento Peak Observatory and have a resolution of 1.5 s of arc. Data include absolute intensities, velocities, linewidths and the optical depth at the center of H. Electron densities range between 6 × 10¹⁰cm^–3 and 1.2 × 10¹¹cm^–3, while the range of the electron temperature is between 12000K and 15000K at a height of 5400 km above the limb. The analysis of the variation of the line of sight velocity across spicules shows that their rotation is very small. There is very good correlation in the structure of all three lines.     

Extreme ultraviolet spectra of solar active regions and their analysis

Extreme ultraviolet spectra of several active regions are presented and analyzed. Spectral intensities of 3 active regions observed with the NRL Skylab XUV spectroheliograph (170–630 Å) are derived. From this data density sensitive line ratios of Mg viii, Si x, S xii, Fe ix, Fe x, Fe xi, Fe xii, Fe xiii, Fe xiv, and Fe xv are examined and typically yield, to within a factor of 2, electron pressures of 1 dyne cm^–2 (n _e T = 6 × 10¹⁵ cm^–3 K). The differential emission measure of the brightest 35 × 35 portion of an active region is obtained between 1.4 × 10⁴ K and 5 × 10⁶ K from HCO OSO-VI XUV (280–1370 Å) spectra published by Dupree et al. (1973). Stigmatic EUV spectra (1170–1710 Å) obtained by the NRL High Resolution Telescope and Spectrograph (HRTS) are also presented. Doppler velocities as a function of position along the slit are derived in an active region plage and sunspot. The velocities are based on an absolute wavelength scale derived from neutral chromospheric lines and are accurate to ±2 km s^–1. Downflows at 10⁵ K are found throughout the plage with typical velocities of 10 km s^–1. In the sunspot, downflows are typically 5 to 20 km s^–1 over the umbra and zero over the penumbra. In addition localized 90 and 150 km s^–1 downflows are found in the umbra in the same 1 × 1 resolution elements which contain the lower velocity downflows. Spectral intensities and velocities in a typical plage 1 resolution element are derived. The velocities are greatest ( 10 km s^–1) at 10⁵ K with lower velocities at higher and lower temperatures. The differential emission measure between 1.3 × 10⁴ K and 2 × 10⁶ K is derived and is found to be comparable to that derived from the OSO-VI data. An electron pressure of 1.4 dynes cm^–2 (n _e T = 1.0 × 10¹⁶ cm^–3 K) is determined from pressure sensitive line ratios of Si iii, O iv, and N iv. From the data presented it is shown that convection plays a major role in determining the structure and dynamics of the active region transition zone and corona.     

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.     

Evidence for supernova-induced star formation in the monoceros OB2 association

The Monoceros ring, a circular optical nebulosity 3°.5 in diameter and centred at R.A.=6^h37^m, Dec.=6°30 (l ⁱⁱ =205°.5,b ⁱⁱ =0°.2) is in good structural agreement with radio observations. A neutral hydrogen shell is also accurately projected on the ring. These observations are consistent with the Monoceros ring being a supernova remnant 90–100 pc in diameter expanding at about 45 km s^–1 and having an age of the order of a million years. Bright Hii regions containing early-type stars (e.g., galactic cluster NGC 2244 in the Rosette nebula) and extremely young stars of the OB association Mon OB2 lie at the edges of the ring. The positional and temporal coincidence of the Mon OB2 association with a supernova remnant suggests that probably the star formation in this region is induced or speeded up by the passage of a supernova shock wave through the clumpy interstellar medium.     

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.     

Structure of the expanding atmosphere of P Cygni

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).