Oscillations in the Umbral Atmosphere

The results of simultaneous observations of oscillations in the chromosphere, transition region, and corona above nine sunspots are presented. The data are obtained through coordinated observing with the Solar and Heliospheric Observatory — SOHO and the Transition Region And Coronal Explorer — TRACE. Oscillations are detected above each umbra. The power spectra show one dominant frequency corresponding to a period close to 3 min. We show that the oscillations in the sunspot transition region can be modeled by upwardly propagating acoustic waves. In the corona the oscillations are limited to small regions that often coincide with the endpoints of sunspot coronal loops. Spectral observations show that oscillations in the corona contribute to the observed oscillations in the TRACE 171 Å channel observations. We show that a recent suggestion regarding a connection between sunspot plumes and 3-min oscillations conflicts with the observations.     

Oscillations and waves in a sunspot

Observations have been made in H of the vertical velocity distribution in a sunspot. Over the umbra the pattern consists of structures of scale-size 2–3. The velocity distribution undergoes oscillations with a period of about 165 s and typical amplitude ±3 km s^–1, but the pattern breaks down after one or two cycles because the period of oscillation varies typically by ±20 s from place to place. Transverse waves develop in the outer 0.1 of the umbral radius and propagate outwards with a velocity of about 20 km s^–1, becoming gradually invisible by or before the outer penumbral boundary; the amplitude is about ±1 km s^–1 at the umbra-penumbra border.The penumbral waves are believed to be basically of the Alfvén type, with 3 × 10^–8 g cm^–3. The umbral oscillations presumably represent gravity waves. In both cases the fluxes are inadequate by two orders of magnitude to account for the sunspot energy deficit.     

Sunspot Plumes and Flow Channels

It is well known that sunspots are dark. This statement is not correct in the sunspot atmosphere between the chromosphere and the corona, where sunspots often are brighter than their surroundings. The brightest feature in the sunspot transition region is called a sunspot plume. Not all sunspots contain a plume. We find that 20 out of 21 sunspots show a plume when one magnetic polarity dominates the sunspot region out to a distance of 50 from the sunspot. Most sunspots show downflows that exceed 25 km s^–1 in the sunspot plumes at temperatures close to 250000 K. This downflow is not maintained by inflow from the corona, but by gas at transition region temperatures, streaming in flow channels from locations well outside the sunspot. We suggest that this inflow is a necessary requirement for the sunspot plume to occur and present a working hypothesis for the origin of sunspot plumes. This paper is the first thorough spectral analysis of sunspot plumes. It is based on simultaneous observations of ten or six EUV emission lines in 42 sunspot regions with the Coronal Diagnostic Spectrometer – CDS on the Solar and Heliospheric Observatory – SOHO. The line profiles are studied in detail with another SOHO instrument, the Solar Ultraviolet Measurements of Emitted Radiation – SUMER.     

Kinematics of solar prominences

Two sequences of OSO-4 spectroheliograms in Mg x and Si xii obtained during October–November 1967 and covering the intervals of 83 and 22 hr, respectively, have been analyzed to reveal quasi-periodic oscillations of EUV flux from solar sources with a periodicity of 5–14 hr. The oscillation periods of the emission flux from local sources over sunspots and magnetic field enhancements in plages without spots have been investigated in correlation with characteristics of the respective AR and plages. The greatest periods (> 8 hr) are shown to be peculiar of small sunspots or sunspot groups at the initial or final stage of their development, whereas the smallest periods ( 5–6 hr) are observed in the case of large well-developed groups at the maximum stage of development. In quiet regions on the Sun and plages without spots, the oscillation periods are 6–8 hr. The surface areas in which the oscillations are synchronous and coincide in phase have typical dimensions of 1 in quiet and 1 to 5 in active regions. These areas form a spatial structure similar to the chromospheric network and supergranules. The characteristic lifetime of the structure elements is 1.5–2 days.     

Umbral oscillations and penumbral waves in Hα

We present examples of umbral oscillations observed on Big Bear H filtergram movies and investigate the relation between umbral oscillations and running penumbral waves occurring in the same sunspot. Umbral oscillations near the center of the umbra are probably physically independent of the penumbral waves because the period of these umbral oscillations (150 s) is shorter than the penumbral wave period (270 s) but not a harmonic. We also report dark puffs which emerge from the edge of the umbra and move outward across the penumbra, and which have the same period as the running penumbral waves. We interpret these dark puffs to be the extension of chromospheric umbral oscillations at the edge of the umbra. It is suggested that the dark puffs and the running penumbral waves have a common source: photospheric oscillations just inside the umbra.     

Short-Period Waves That Heat the Corona Detected at the 1999 Eclipse

As a part of a study of the cause of solar coronal heating, we searched for high-frequency (1 Hz) intensity oscillations in coronal loops in the [Fexiv] coronal green line. We summarize results from observations made at the 11 August 1999 total solar eclipse from Râmnicu-Vâlcea, Romania, through clear skies. We discuss the image reduction and analysis through two simultaneous series of coronal CCD images digitized at 10 Hz for a total time of about 140 s. One series of images was taken through a 3.6 Å filter isolating the 5303 Å[Fexiv] coronal green line and the other through a 100 Å filter in the nearby K-corona continuum. Previous observations, described in Pasachoff et al. (2000), showed no evidence for oscillations in the [Fexiv] green line at a level greater than 2% of coronal intensity. We describe several improvements made over the 1998 eclipse that led to increased image clarity and sensitivity. The corona was brighter in 1999 with the solar maximum, further improving the data. We use Fourier analysis to search in the [Fexiv] channel for intensity oscillations in loops at the base of the corona. Such oscillations in the 1-Hz range are predicted as a result of density fluctuations from the resonant absorption of MHD waves. The dissipation of a significant amount of mechanical energy from the photosphere into the corona through this mechanism could provide sufficient energy to heat the corona. A Monte Carlo model of the data suggests the presence of enhanced power, particularly in the 0.75–1.0 Hz range, and we conclude that MHD waves remain a viable method for coronal heating.     

Solar cycle variation of sunspot intensity

Broad band pinhole photometer intensity observations of 15 large sunspots covering the spectral region 0.387–2.35 m are presented. The data are based on measurements on approximately 500 days during the period June, 1967 to December, 1979.We have found real and significant intensity differences between large sunspots. These differences may be explained by a systematic variation in the umbral temperature throughout the solar cycle. A connection between umbra intensity and heliographic latitude is discussed.No center-limb variation in the umbra/photosphere intensity ratio is detected. We have searched for possible connections between umbra intensity and a number of other sunspot parameters, like the spot size, without detecting any significant correlation. We conclude that the umbra/photosphere intensity ratio seems to be a unique function of epoch for large sunspots.     

Photospheric and chromospheric oscillations in the base of coronal holes

In this paper we carry out an analysis of the spatial–temporal line-of-sight velocity variations measured in the chromospheric (H, H) and photospheric (Fei 6569 Å, Fei 4864 Å, Nii 4857 Å) lines at the base of 17 coronal holes. Time series of a duration from 43 to 120 min were recorded with the CCD line-array and the CCD matrix. Rather frequently we observed quasi-stationary upward flows with a measured velocity of up to 1 km s^–1 in the photosphere and up to 4–5 km s^–1 in the chromosphere (equivalent radial velocity of up to 3 km s^–1 and up to 12–15 km s^–1 accordingly) near dark points on the chromospheric network boundary inside polar CH. Line-of-sight velocity fluctuation spectra contain meaningful maxima in the low-frequency region clustering around the values 0.4, 0.75, and 1 mHz. Usually, the spatial localization of these maxima mutually coincides and, in our opinion, coincides with the chromospheric network boundary. Acoustic 3- and 5-min oscillations are enhanced in the coronal hole region and reach 1 km s^–1 in the photosphere and 3–4 km s^–1 in the chromosphere. These oscillations are not localized spatially and are distinguished throughout the entire region observed.     

Brightness variations of the white light corona during the years 1964–67

Observations of the white light corona were made on over 900 days during the years 1964–67 at heights between 1.125 and 2.0 R with the K-coronameter at Mount Haleakala and Mauna Loa, Hawaii. The brightness distribution of the minimum corona was elliptical with average equatorial intensities three times the polar. Coronal features of the new cycle at 1.125 R occurred predominantly in the sunspot zones at 25–30° latitude and in a high latitude zone which migrated toward the North pole before solar maximum. The brightness of the inner corona doubled over this period and a close association is found between the average corona and 10.7-cm solar radio flux. Electron densities in the equatorial regions were nearly twice those of Van de Hulst's model corona, in agreement with the results of recent eclipse observations.At Hawaii Institute of Geophysics.     

Imaging Spectropolarimetry of Ti I 2231 nm in a Sunspot

Spectro-polarimetric observations at 2231 nm were made of NOAA 10008 near the west solar limb on 29 June 2002 using the National Solar Observatory McMath–Pierce Telescope at Kitt Peak and the California State University Northridge – National Solar Observatory infrared camera. Scans of spectra in both Stokes I and Stokes V were collected; the intensity spectra were processed to remove strong telluric absorption lines, and the Stokes V umbral spectra were corrected for instrumental polarization. The sunspot temperature is computed using the continuum contrast and umbral temperatures down to about 3700 K are observed. A strong Tii line at 2231.0 nm is used to probe the magnetic and velocity fields in the spot umbra and penumbra. Measurements of the Tii equivalent width versus plasma temperature in the sunspot agree with model predictions. Zeeman splitting measurements of the Stokes I and Stokes V profiles show magnetic fields up to 3300 G in the umbra, and a dependence of the magnetic field on the plasma temperature similar to that which was seen using Fei 1565 nm observations of the same spot two days earlier. The umbral Doppler velocity measurements are averaged in 16 azimuthal bins, and no radial flows are revealed to a limit of ±200 m s^–1. A Stokes V magnetogram shows a reversal of the line-of-sight magnetic component between the limb and disk center sides of the penumbra. Because the Tii line is weak in the penumbra, individual spectra are averaged in azimuthal bins over the entire penumbral radial extent. The averaged Stokes V spectra show a magnetic reversal as a function of sunspot azimuthal angle. The mean penumbral magnetic field as measured with the Stokes V Zeeman component splitting is 1400 G. Several weak spectral lines are observed in the sunspot and the variation of the equivalent width versus temperature for four lines is examined. If these lines are from molecules, it is possible that lines at 2230.67, 2230.77, and 2231.70 nm originate from OH, while the line at 2232.21 nm may originate from CN.     

Broad-band circular polarimetry of sunspots, 0.4–1.7 microns: Spatial scans with a 3.4 arc sec diameter aperture

Spatial scans with a resolution of 3.4 arc sec of the broad-band circular polarization of several sunspots have been made in five filter bands over the wavelength range 0.4–1.7µ with a sensitivity of 1 × 10^–6 fractional polarization. The scans, across a spot through the penumbra and umbra center, revealed two important features: (1) The broad-band circularly polarized fluxV reverses in sign, or diminishes to near zero, at the center of the umbral region relative to the outer penumbra. This effect was wavelength dependent and was most clearly detected as a definite reversal in a band at 1.2µ, although a reversal was also detected in a very broad band extending from 0.8 to 1.6µ. (2) There is a marked asymmetry: in all cases the limbward penumbral region exhibited strongerV values than did the disk-center (inward) side of the spot, at all observed wavelengths. Such previously unreported structure in the magnetic circular polarization of sunspots provides new clues for understanding the anomalous large broad-band polarization at short wavelengths and at the same time imposes new constraints on sunspot models. For example, the polarization reversal in the umbra relative to the penumbra can be naively explained by return-flux sunspot models; but this is not the only interpretation. Alternatively, it can relate to reversals in mass-flow velocities and/or vertical velocity gradients, as between the umbra and penumbra.     

The oscillatory velocity field observed in a unipolar sunspot region

The velocity field has been mapped for 42 min in an area 80 by 85 containing a unipolar sunspot. Apparent shifts of Fe i 5233 were measured photoelectrically using a rectangular scanning aperture 1.6 × 4.0. The sunspot did not exert a marked influence on the generally random pattern of oscillations at a period of 300 s. Discrete periods of oscillation both longer and shorter than 300 s were excited within the enhanced magnetic field boundaries of this spot. Umbral oscillations at periods near 180 s were detected in agreement with independent observations of the same spot during the previous solar rotation.NRC Postdoctoral Fellow, 1969–71.     

Orientation of the galactic magnetic field in the vicinity of the solar system

It is found from analysis of the position angles of the plane of polarization of about 3000 stars (¦b¦ 5° andP 0.5%) that the angle between the magnetic field and the equatorial plane of the galaxy is approximately 0–5°. The distance within which the local magnetic fields of the galaxy have a greater effect on the position angles of the plane of polarization than the galactic magnetic field is estimated to be about 500 pc. The effect of the galactic magnetic field becomes dominant for distancesr 1000 pc.Translated fromAstrofizika, Vol. 39, No. 4, pp. 553–559, November, 1996.     

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.     

On the sunspot transition region

The EUV line emission and relative line-of-sight velocity in the transition region between the chromosphere and corona of 36 sunspot regions are investigated, based on observations with the Coronal Diagnostic Spectrometer – CDS and the Solar Ultraviolet Measurements of Emitted Radiation – SUMER on the Solar and Heliospheric Observatory – SOHO. The most prominent features in the transition-region intensity maps are the sunspot plumes. In the temperature range between log T=5.2 and log T=5.6 we find that 29 of the 36 sunspots contain one or two sunspot plumes. The relative line-of-sight velocity in sunspot plumes is high and directed into the Sun in the transition region, for 19 of the sunspots the maximum velocity exceeds 25 km s^?1. The velocity increases with increasing temperature, reaches a maximum close to log T=5.5 and then decreases abruptly.

Attention is given to the properties of oscillations with a period of 3 min in the sunspot transition region, based on observations of six sunspots. Comparing loci with the same phase we find that the 3-min oscillations affect the entire umbral transition region and part of the penumbral transition region. Above the umbra the observed relation between the oscillations in peak line intensity and line-of-sight velocity is compatible with the hypothesis that the oscillations are caused by upward-propagating acoustic waves. Information about intensity oscillations in the low corona is obtained from observations of one sunspot in the 171 Å channel with the Transition Region And Coronal Explorer – TRACE. We conclude that we observe the 3-min sunspot oscillations in the chromosphere, the transition region and the low corona. The oscillations are observable over a wider temperature range than the sunspot plumes, and show a different spatial distribution than that of the plumes.

     

Did Predictions of the Maximum Sunspot Number for Solar Cycle 23 Come True?

For solar cycle 23, the maximum sunspot number was predicted by several workers, and the range was very wide, 80–210. Cycle 23 started in 1996 and seems to have peaked in 2000, with a smoothed sunspot number maximum of 122. From about 20 predictions, 8 were within 122±20. There is an indication that a long-term oscillation of 80–100 years may be operative and might have peaked near cycle 20 (1970), and sunspot maxima in cycles in the near future may be smaller and smaller for the next 50 years or so and rebound thereafter in the next 50 years or so.     

Cluster analysis of the space-time distribution of sunspot groups during solar cycle no. 20

Cluster analysis (a Bayesian iteration procedure) was used to study the space-time distribution of sunspot groups in the time interval from 1965 to 1977. (Data were taken from the Greenwich and Debrecen Heliographic Results.) The distribution proved to be significantly non-random for the 8–10 groups cluster^–1 (gr cl^–1) level of clustering. Convincing evidence also favours non-random behaviour for other levels of clustering from the lowest (3–4 gr cl^–1) up to the highest ( 150 gr cl^–1) level. The rotation rate of the non-random pattern is generally slightly lower than the Carrington rate.The 8–10 gr cl^–1 level, crudely corresponding to the sunspot nests investigated earlier, was studied in more detail. The cycle- and latitude-averaged rotational rate of the nests is slightly ( 1%) but significantly lower than the Carrington rate. Their differential rotation is strongly reduced: the cycle-averaged rotational rate varies only by 2–3% within the sunspot belt. A slight but significant bimodality is seen in the differential rotation curve: the intermediate latitudes ( 10°–20°) show a somewhat slower rotation than both the equatorial and the higher latitude regions. This might be explained by a time-dependence of the rotation rate coupled with the butterfly diagram.     

Euv Spectroscopy of the Sunspot Region Noaa 7981 Using Soho – I. Line Emission and Time Dependence

EUV spectra of a medium-size sunspot and its surroundings, NOAA 7981, were obtained on 2 August 1996 with the Coronal Diagnostic Spectrometer (CDS) and the Solar Ultraviolet Measurements of Emitted Radiation (SUMER) on the Solar and Heliospheric Observatory (SOHO). The spectral lines formed in the transition region and corona show considerable structure and large deviations from a uniform spatial distribution over the active region. Enhanced EUV emissions in transition region lines are concentrated in small regions outside the umbra of the sunspot throughout most of the observing sequence. Only during a short, active period do we find an enhanced line emission that reaches into the umbra. Preliminary values for the umbral intensity are given.Marked changes are detected between the spatial distribution of line emission in the transition region and the low corona. The difference is not limited to cool and hot non-flaring loops not being cospatial, but includes differences both regarding the time variability and the orientation and size of the emission features. Whereas both rapid ( 4 and 2 min) and slow ( 10 and 12 h) temporal variations are found in the chromosphere (Hei 584 Ú) and transition region (Ov 629 Ú), the response in the low corona (Mgix 368 Ú) is slow ( 5 h). Furthermore, marked differences between the spatial distributions in the Mgviii 315 Ú, Mgix 368 Ú lines formed in the low corona and the coronal Fexiv 334 Ú, Fexvi 360 Ú lines are detected.     

A Molecular Ring in the Liner Ngc 5218

High resolution OVRO CO 1–0 observations of the inner 30 in the LINER galaxy NGC 5218 reveal the presence of a double centrally peaked molecular concentration with extensions out to a radius of 12. The molecular mass detected is 2.4 × 10⁹ M and the gas surface density is high, 3000 M pc square in the inner 500 pc. The SFR is 2–3 M yr^–1 and the SFE is 13, which are low or moderate values for that gas surface density. We interpret the inner feature as a rotating molecular ring with a radius of 200 pc. We furthermore suggest that the LINER activity in NGC 5218 is not caused by an aging starburst, but by a buried AGN.     

Pulsars: Observational parameters and a discussion on dispersion measures

The relevant data for the known 147 pulsars are presented in graphical and tabular forms. Various data correlations are discussed, and a detailed analysis of pulsar dispersion measures and distances is given. The range of the electron densities in the diffuse interstellar medium is found to be 0.01 cm^–3n _e0.1 cm^–3, and n _e0.03 cm^–3. The dispersion scale height for pulsars is found to be 5.9±0.7 pc cm^–3 implying a linear scale height of 200 pc, which is much smaller than the electron scale height of our Galaxy.Astrophysics and Space Science Review Paper.