Some properties of velocity fields in the solar photosphere

Photoelectric measurements of Doppler shifts of various Fraunhofer lines obtained with the Capri magnetograph were analysed. The height dependence of the supergranular and oscillatory motions, as well as the two dimensional structure of these velocity fields is investigated. The most interesting results are the following:

1. The oscillatory and supergranular motions are still clearly present in very deep photospheric layers as detected e.g. by means of the Ci line at 5380.3 Å.
2. Whereas the vertical motions (both of oscillation and supergranulation) increase with height, the horizontal component of the supergranular flow is found to be decreasing slightly.
3. Aperiodic horizontal motions are observed in the photospheric layers, which are probably connected with the process of excitation of the oscillatory field.
4. There is no simple way of describing the oscillatory field in terms of independently oscillating ‘cells’, since the two-dimensional pattern changes its appearance drastically already in a fraction of one oscillation period.
5. The correlation obtained by previous observers between vertical stationary motions, the chromospheric network and magnetic fields in particular is confirmed.

     

Wave systems in the chromosphere

Observations are reported of two, possibly three, distinct wave systems in the Hα chromosphere.

1. Velocity films show waves propagating predominantly outwards along mottles and fibrils from as close as 2000 km to the network axis at velocities of the order of 70 km s^-1. The line-of-sight component of the velocity amplitude is estimated to be typically 5 km s^-1. The velocities are accompanied by propagating intensity fluctuations. The system is interpreted as one of basically Alfvén waves. Similar waves are observed propagating predominantly outwards along superpenumbral fibrils radiating from a small sunspot.
2. The velocities in the chromospheric granulation undergo fluctuations of an oscillatory character but without any observable horizontal propagation. The intensities show a close correlation with the velocities, maximum intensity occurring about T/4 after maximum downward velocity. The period is variable across the surface (2.5 min upwards). The intensity-velocity correlation is characteristic of a standing compressional wave.
3. Intensity cinefilms at Hα line centre show in places a horizontal drift of the chromospheric granulation pattern at about 12 km s^-1 without any accompanying vertical velocity fluctuations. It is not known whether this is due to a gas stream at sonic velocities, or to a horizontally propagating sound wave.

The Alfvén wave system is shown to make a significant contribution to coronal heating. Whether the velocity fluctuations in the chromospheric granulation also make an important contribution depends on whether there are upwardly propagating or standing waves; this is not yet established despite the intensity-velocity correlation.     

The photospheric network

Spectroheliograms, obtained in certain Fraunhofer lines with the 82-cm solar image at the Kitt Peak National Observatory, show a bright photospheric network having the following properties:

1. It resembles, but does not coincide with, the chromospheric network, the structure of the photospheric network being finer and more delicate than the relatively coarse structure of the chromospheric network.
2. It is exactly cospatial with the network of non-sunspot photospheric magnetic fields.
3. Its visibility in a given photospheric Fraunhofer line is primarily dependent on the states of ionization and excitation from which the line is formed and secondarily dependent on the Zeemansensitivity of the line-being most visible in low-excitation lines of neutral atoms and least visible in high-excitation lines of singly ionized atoms.

We conclude that these magnetic regions of the solar atmosphere are a few hundred degrees hotter than their surroundings, and that they are visible in white light near the limb as photospheric faculae.     

Spectral analyses of solar photospheric fluctuations

Successful subtraction of instrumental background variations has permitted spectral analyses of two-dimensional measurement arrays of granulation brightness fluctuations at the center of the disk, arrays obtained from Stratoscope I, 1959B-flight, high-resolution frames B1551 and B3241.

1. RMS's, uncorrected for instrumental blurring, are 0.0850 of mean intensity for B1551 and 0.0736 for B3241, somewhat higher than other determinations. These between-frame and between-investigation differences probably result from a combination of calibration errors, frame resolution differences, and, most likely, granulation pattern differences.
2. Significant variations over each array of mean intensities and RMS's, determined for sub-arrays with dimensions in the 2500–10000 km range, indicate spatial brightness and RMS variations larger than the ‘scale’ of the granulation pattern, supporting a turbulent interpretation of photospheric convection.
3. One-dimensional power-spectra shapes provide objective and discriminating criteria for determining granulation pattern differences and, possibly, frame resolution.
4. Two-dimensional power spectra show small, essentially random deviations from axial symmetry which lie almost entirely within the 50% confidence limits.
5. Spectral densities and fluctuation power spectra, computed from the two-dimensional power spectra and corrected for instrumental blurring, noise, and blemishes, have a useable radial wavenumber range nearly double that of earlier Stratoscope I analyses.
6. Corrected RMS's obtained from the corrected fluctuation power spectra, 0.145 ± 0.046 for B1551 and 0.136 ± 0.048 for B3241, depend critically on the accuracy of the correction.
7. The spectra's wavenumber range includes the granulation-fluctuation-producing domain but not the Kolmogoroff domain of turbulence spectra.

     

Small-scale motions over concentrated magnetic regions of the quiet Sun

We have used a 5.5 min time-sequence of spectra in the Fe i lines λ5576 (magnetically insensitive), λ6301.5 and λ6302.5 (magnetically sensitive) to study the association of concentrated magnetic regions and velocity in the quiet Sun. After the elimination of photospheric oscillations we found downflows of 100–300 m s ^?1, displaced by about 2″ from the peaks of the magnetic field; this velocity is comparable to downflow velocity associated with the granulation and of the same order or smaller than the oscillation amplitude. Quasi-periodic time variations of the vertical component of the magnetic field up to ± 40% were also found with a period near 250 s, close to the values found for the velocity field. Finally we report a possible association of intensity maxima at the line center with peaks of the oscillation amplitude.     

Direct Measurement Results of the Time Lag of LOS-Velocity Oscillations Between Two Heights in Solar Faculae and Sunspots

We present an investigation of line-of-sight (LOS) velocity oscillations in solar faculae and sunspots. To study the phase relations between chromospheric and photospheric oscillations of the LOS velocity, we measured the time lag of the chromospheric signal relative to the photospheric one for several faculae and sunspots in a set of spectral line pairs. The measured time lags are different for different objects. The mean measured delay between the oscillations in the five-minute band in faculae is 50?s for the Si?i 10?827?Å?–?He?i 10?830?Å pair; for the pair Fe?i 6569?Å?–?Hα 6563?Å the mean delay is 20?s; for the pair Fe?i 4551?Å?–?Ba?ii 4554?Å the mean delay is 7?s; for the pair Si?i 8536?Å?–?Ca?ii 8542?Å the mean delay is 20?s. For the oscillations in the three-minute band in sunspot umbrae the mean delay is 55?s for the Si?i 10?827?Å?–?He?i 10?830?Å pair; for the Fe?i 6569?Å?–?Hα 6563?Å pair it was not possible to determine the delay; for the Fe?i 4551?Å?–?Ba?ii 4554?Å pair the mean delay is 6?s; for the Si?i 8536?Å?–?Ca?ii 8542?Å pair the mean delay is 21?s. Measured delays correspond to the wave propagation speed, which significantly exceeds the generally adopted speed of sound in the photosphere. This raises the question of the origin of these oscillations. The possibility that we deal with slow MHD waves is not ruled out.     

Étude morphologique et cinématique des structures fines d'une tache solaire

A sequence of 34 photographs of the main spot of the group H 26 (Daily Maps of the Sun, Freiburg 1970, Rome number 5847) has been obtained with the 38 cm refractor of the Pic-du-Midi Observatory, showing throughout a resolution very close or equal to 0′'.3. An interval of 3 hr is covered. The pictures taken at intervals of 6 min approximately permit to study the fine structure of the penumbra and associated phenomena:

1. The penumbra appears to consist of bright grains, generally lined up in the form of filaments, showing up against a dark background (see Figure 1).
2. The bright grains form all over the penumbra (see Figure 5).
3. They move toward the umbra of the spot. Their horizontal velocity is zero at the border penumbra-photosphere and maximum at the umbral border (0.5 km s^?1) (see Figures 3,4 and 8). Therefore, the grains never originate in the photosphere nor do they enter it.
4. They disappear in the penumbra proper or, if they form near enough to the umbra and live long enough, they can enter the umbra and their appearance becomes similar to that of umbral dots.
5. The life time of the grains is a function of their place of origin within the penumbra: It is maximum and of the order of 3 hr or more for those forming in the middle part of the penumbra, and 50 and 40 min respectively for the points formed in the inner and outer part of the penumbra.

     

The evolution of the photospheric network

A time-lapse sequence of spectroheliograms in the bandhead of CN at λ3883 reveals the following behavior of the photospheric network with time:

1. There is a steady flow of bright ‘points’ (? 1000 km in diameter) laterally outward from sunspots at speeds on the order of 1 km·sec^?1. After traveling about 10 000 km from a sunspot they either conglomerate to form fragments of the photospheric network or disappear.
2. Spatial changes in the network pattern seem to take place by means of the shifting of network fragments laterally on the solar surface. Although most small-scale details are recognizable after 5–10 minutes, within 30 minutes nearly all the details have changed completely. In contrast to this, the large-scale network pattern seems relatively unchanged after 2 1/2 hours.
3. Occasionally ‘new’ network, not resulting from the lateral motion of bright features from either previously existing network or sunspots, appears on the solar surface. This process consists of the formation in approximately 10 minutes of bright points and a darker-than-average feature between them. The dark feature disappears in another 5–10 minutes and the bright points separate at a relative speed of a few km·sec^?1. If the event is of a sufficiently large magnitude, a sunspot will appear.

These observed changes of the photospheric network with time are interpreted as formation and motions of photospheric magnetic fields. It is suggested that these motions reflect the presence of both short-lived small-scale and long-lived large-scale photospheric currents such as one might expect from the granulation and the supergranulation.     

Spectral observations of spicules at two heights in the solar chromosphere

An observational program at the Sacramento Peak Observatory in 1965 provided high-dispersion spectra of the solar chromosphere in several spectral regions simultaneously. These regions included various combinations of the spectral lines Hα, Hβ and H?, the D₃-line of Hei, the infrared triplet of Oi, and the H- and K-lines and the infrared triplet of Caii. With the use of an image slicer the observations were made simultaneously at two heights in the solar chromosphere separated by several thousand kilometers. From these data we draw the following conclusions:

1. Emission of different lines arises in the same chromospheric features. The intensity ratio of lines of different elements varies significantly from spicule to spicule. For the H- and K-lines of ionized calcium, this ratio remains constant, independent of wavelength throughout the line, overall intensity, and height in the chromosphere. Two rare-earth lines in the wing of the H-line show no spicular structure at all.
2. The line-of-sight velocities of many features reverse as a function of time, although most spicules show velocities in only one direction. The simultaneous spectra at two heights show most spicules to have the same line-of-sight velocity at both. There may be an additional class of features, mostly rapidly moving, whose members have line-of-sight velocities that increase with height. These features comprise perhaps 10% of the total. Velocity changes occur simultaneously, to within 20 sec, at two heights separated by 1800 km, indicating velocities of propagation of hundreds of km/sec. The velocity field of individual features is often quite complicated; many spectral features are inclined to the direction of dispersion, implying that differential mass motions are present.
3. The existence of anomalously broad H and K profiles is real. Even with high dispersion and the best seeing, such profiles are not resolved into smaller features. The central reversal in K, H and Hα appears to remain unshifted when the wings are displaced in wavelength, indicating that the reversal is non-spicular.

     

Orbital eccentricity study for the spherical component of our galaxy

An analysis of the data concerning high-velocity stars from Eggen's catalogue aimed at a determination of the approximate slope of the mass function for the spherical component of our Galaxy, and at estimating the local circular velocity, as well as the local rotation velocity, as by-products, has been performed. Our conclusions are that:

1. A linear dependence of the mass on the radius is very likely;
2. the value of the limiting radius is most likely equal to (40±10) kpc;
3. the two local velocities are approximately equal to each other, being both equal to (230±30) km s^?1;
4. the local escape velocity appears to be most likely equal to (520±30) km s^?1;
5. the total mass of a corona, obtained in this way, is (5±1)×10¹¹ M _⊙.

     

A comparison of He ii 304 Å and He i 10 830 Å spectroheliograms

Spectroheliograms were obtained simultaneously in the He ii 304 Å emission line and the He i 10 830 Å absorption line with an angular resolution of approximately 5″. A negative print of the 304 Å image is matched with a positive print of the 10 830 Å image so that corresponding features of the chromospheric network (including active regions) appear identical in the two images. Differences between these images include the facts that:

1. Disk filaments and limb darkening are strongly visible in the 10 830 Å positive image, but they are weakly visible (as lightenings) in the 304 Å negative image.
2. The contrast between the chromospheric network and the network cell centers is much greater in the 10 830 Å image than in the 304 Å negative image.

These results provide constraints on models of helium line formation in various types of solar features.     

Drift theory of charged particles in electric and magnetic fields

In this paper we review the drift theory of charged particles in electric and magnetic fields. No new physical interpretations are added to this classical topic, but through an alternative, simplified derivation of the guiding centre velocity, several complexities are eliminated and possible misconceptions of the theory are clarified. It is shown that:

1. The curvature/gradient drift velocity in the magnetic field, averaged over a particle distribution function is to lowest order in the direction of?×B/B ², while the average particle velocity is in the direction ofB×? P withP the scalar particle pressure.
2. These drift directions are correct for first-order expansions of the particle distribution function, and only second-order or higher expansions change these directions.
3. The?×B/B ² drift, which is the standard gradient plus curvature drift, and which is usually considered as a ‘single particle’ drift, need not be ‘reconciled’ with theB×? P, or ‘macroscopic, collective’ drift, as is often asserted in the literature. They are in fact related per definition and we show how.
4. When viewed in fixed momentum intervals (p,p+dp), the so-called Compton-Getting factor enters into the electric field (E×B)/B ² drift term.
5. The results are independent of the scale length of variation ofE andB, in contrast to existing drift theory. We discuss the implications of this result for three important cases.

     

Optical observations of the interstellar absorption lines towards the M8 Nebula

High resolution optical observations (FWHM ～ 10–13 km s^?1) of the I-S gas towards the early-type stars HD 164 794, HD 164816, and HD 165052 in the M8 Nebula are presented. A high velocity componentV _LSR=?26 km s^?1 has been detected in all 3 stars' spectra. A line profile fitting analysis has been carried out on the observed Caii and Nai absorption lines to determine cloud component column densities and to subsequently determine the physical and chemical conditions of the associated I-S gas.     

An evaluation of the possibility of studying flare plasma turbulence using the satellites of Hei line forbidden components

Using the Baranger-Mozer method, we explore the possibility of diagnosing the flare plasma of forbidden Hei lines, that permits the determination of the plasma oscillation frequency and noise level. Examination of the Hei lines observed in solar flare has led us to conclude that:

1. the appearance of satellites of forbidden components in the flares spectrum, due to turbulent electric fields, is the most probable for Hei 3819.606 Å lines;
2. the Baranger-Mozer method is more sensitive to the high-frequency component of turbulent fields than to the low-frequency ones;
3. the upper limit of the turbulent oscillation level in flares is evaluated.

In the spectrum of the solar flare of 26 September, 1963 we detected satellites of the forbidden component of the 3820 Å line and used its relative intensity to derive the level of low-frequency oscillations (～1.5 kVcm^-1).     

Optical spectral decomposition of NGC 4051

Considering the host galaxy contribution, a spectral decomposition method is used to reanalyzed the archive data of optical spectra for a narrow line Seyfert 1 galaxy, NGC 4051. The light curves of the continuum f _λ (5100 Å), and Hβ, He ii, Fe ii emission lines are given. We find strong flux correlations between line emissions of Hβ, He ii, Fe ii and the continuum f _λ (5100 Å). These low-ionization lines (Hβ, Fe ii, He ii) have “inverse” intrinsic Baldwin effects. Using the methods of the cross-correlation function and the Monte Carlo simulation, we find the time delays, with respect to the continuum, are $3.45^{+12.0}_{-0.5}~\mbox{days}$ with the probability of 34 % for the intermediate component of Hβ, $6.45^{+13.0}_{-1.0}~\mbox{days}$ with the probability of 65 % for the intermediate component of He ii. From these intermediate components of Hβ and He ii, the calculated central black hole masses are $0.86^{+4.35}_{-0.33}\times 10^{6}$ and $0.82^{+3.12}_{-0.45}\times 10^{6}~M_{\odot }$ . We also find that the time delays for Fe ii are $9.7^{+3.0}_{-5.0}~\mbox{days}$ with the probability of 36 %, $8.45^{+1.0}_{-2.0}~\mbox{days}$ with the probability of 18 % for the total epochs and “subset 1” data, respectively. It seems that the Fe ii emission region is outside of the Hβ emission region.     

Energy changes of cosmic rays in the interplanetary region

A clarification and discussion of the energy changes experienced by cosmic rays in the interplanetary region is presented. It is shown that the mean time rate of change of momentum of cosmic rays reckoned for a fixed volume in a reference frame fixed in the solar system is 〈p〉 =p V·G/3 (p=momentum,V is the solar wind velocity andG=cosmic-ray density gradient). This result is obtained in three ways:

1. by a rearrangement and reinterpretation of the cosmic-ray continuity equation;
2. by using a scattering analysis based on that of Gleeson and Axford (1967);
3. by using a special scattering model in which cosmic-rays are trapped in ‘magnetic boxes’ moving with the solar wind.

The third method also gives the rate of change of momentum of particles within a moving ‘magnetic box’ as 〈p〉_ad = ?p ?·V/3, which is the adiabatic deceleration rate of Parker (1965). We conclude that ‘turnaround’ energy change effects previously considered separately are already included in the equation of transport for cosmic rays.     

A statistical analysis of NaI D1 profile fluctuations at the center of the solar disk

Three radial-velocity fluctuation arrays V(Δλ, Y) and line-formation fluctuation arrays L(Δλ, Y),where Δλ is wavelength displacement from the center of Nai D₁ and Y is displacement on the Sun's surface along the spectrograph slit, were obtained from Sacramento Peak Observatory spectrograms. The variations of these line profile fluctuations are qualitatively described. The RMS_υ's, coherences, and power spectra shapes for V(Δλ, Y) fluctuations are examined at different Δλ with the corresponding effective heights of formation calculated with Mein weighting functions. Results include: (a) possible anticorrelation between continuum fluctuations and those near line center; (b) RMS _υ ^(cr) 's, which are root-mean-square values of the radial velocity corrected for instrumental and atmospheric blurring, are large (1.5 to 4.0 km s^?1) primarily due to large corrections for atmospheric blurring; (c) RMS _υ ^(cr) minima at effective heights of formation above 350 km suggest penetration of granulation velocities into the upper photosphere; (d) very rough determinations of RMS _υ ^(cr) 's, which are additionally corrected for line-of-sight averaging, range from around 5 km s^?1 in the low chromosphere to a sharp minimum ≤ 0.5 km s^?1 located in the upper photosphere; (e) power spectra shapes reflect decreasing average fluctuation scales above the temperature minimum (possibly high-frequency oscillations) and in the low and middle photosphere (possibly penetration of granulation); and (f) RMS _υ ^(cr) 's and average fluctuation scales suggest changes in the resolvable velocity field occurring near the temperature minimum.     

Wave propagation in the quiet solar chromosphere

In order to precise previous results about wave propagation in the quiet chromosphere (N. Mein and P. Mein, 1976), we study the behaviour of Doppler shifts and intensity fluctuations in 3 lines of Ca ii. We use the same observation as in our previous work, that is to say a sequence of spectra lasting 27 mn, taken at Sacramento Peak Observatory solar tower. Results can be summarized as follows:

1. Phase-lag between intensity fluctuations and dopplershifts is always near 90° in the Ca ii lines, even for frequencies as high as 15 mHz, and whatever is the location in the chromospheric network.
2. Magneto-acoustic waves propagating vertically in a vertical or horizontal magnetic field could account for the observations only if they were, on one hand reflected in the upper atmosphere, on the other hand propagating with a very high sound or Alfvén speed. The lower limit for the speed (70 km s^-1) does not seem to be realistic. Oblique waves could be investigated for better agreement.

     

O VI in the local interstellar medium

We report preliminary results of a search for O VI absorption in the spectra of ～100 hot DA white dwarfs observed by the FUSE satellite. We have carried out a detailed analysis of the radial velocities of interstellar and (where present) stellar absorption lines for the entire sample of stars. In many cases, the velocity differences between the interstellar and photospheric components are below the resolution of the FUSE spectrographs. However, in a significant number of cases the interstellar and photospheric contributions can be separated. In the majority of stars where we find O VI absorption lines, the material is clearly associated with the stellar photosphere and not the interstellar medium. There are a small number of lines-of-sight where the gas is interstellar in nature but the stars are located beyond the boundaries of the local cavity.     

The structure of a sunspot

White-light photographs of a fairly regular sunspot have been obtained for all but one day of its passage across the disk. From microphotometer tracings across these photographs, intensity profiles across the spot have been obtained at several heliocentric angles, θ. Apparent sunspot, umbral and penumbral widths, have been obtained from these profiles, and an examination of these reveals that the well-known Wilson effect (Wilson, 1774) is a rather complex phenomenon comprising four main features:

1. The intensity profiles become increasingly asymmetric at large θ. The penumbra remote from the limb is poorly defined while the penumbral intensity plateau nearer the limb is well defined and sometimes enhanced by an intensity maximum near the umbra-penumbra boundary.
2. A gradual decrease in the apparent width of the disk-side penumbra may occur but this effect is barely significant compared with the rms errors of the observations.
3. The apparent width of the limb-side penumbra is independent of θ for θ < 60° but at larger heliocentric angles it increases sharply and by a significant amount.
4. The apparent umbral diameter also shows no θ-dependence for θ < 60° but beyond this it decreases in an almost complementary manner.

A general model for the structure of a sunspot is put forward which readily explains these results in a qualitative manner but it is emphasised that an adequate analysis of sunspot structure based on these observations requires solutions of the three-dimensional equation of radiative transfer.