Hα Filtergrams Research Articles

Configuration and gradual dynamics of prominence-related X-ray coronal cavities

We have analyzed X-ray images of the solar corona obtained by the S-054 telescope on Skylab, together with Hα filtergrams from the Catania Astrophysical Observatory and EUV and magnetic data, to study the morphology and the evolution of the coronal structures associated with prominences (coronal cavities).

A statistical study of spicule inclinations

The apparent angles of more than 5000 quiet region spicules situated at various position angles around the Sun's limb have been measured on off-band Hα filtergrams taken during the years 1972–1975. The counts were made exclusively in projection on the disk, within 5–25° of the limb. The tendency of the average spicule to lean towards or away from the pole is small at most, and at no latitude exceeds 8°. The spread in angles is 30–35°.

H? observations of the August 12, 1975 type III-RS bursts

We present Hα filtergram observations of a number of the Type III-RS (reverse slope) bursts that occurred on August 12, 1975. Solar radio emission was peculiar on that date in that a large number, and proportion, of the usually rare reverse slope bursts were observed (Tarnstrom and Zehntner, 1975). We show that the radio bursts coincide in time with a homologous set of Hα flares located at the limbward edge of spot group Mt. Wilson 19598. We propose a model in which the reverse slope bursts are the downward branches of U bursts, whose upward branches are hidden behind the coronal density enhancement over the spot group.

On the possible mechanism of formation of emission rim in hydrogen filaments

Hα filtergrams of the chromosphere show an emission rim in many hydrogen filaments. We suppose that formation of this rim is due to photospheric radiation reflected by the filament in the direction of the chromosphere. The calculations show that: (1) the maximum contrast of the rim relative to the undisturbed chromosphere amounts to 1.4; (2) the larger the optical thickness of the filament and the closer to the solar limb it is situated, the brighter and wider is the rim; (3) the rim was not observed in filaments whose heights exceeds 10000km above the chromosphere. These results are in close agreement with observations.

Les bulles chromospheriques

Nous etudions des structures chromospheriques qui ont ete identifiees sur des filtrogrammes pris au bord du disque solaire dans la raie Hα, et que nous appelons ‘bulles’ (figures 1 et 2). Un certain nombre de parametres ont ete mesures (tableau I et figure 4), et une discussion est developpee en ce qui concerne les hauteurs maximum et les diametres des bulles. Nous comparons les bulles avec d'autres structures chromospheriques (tableau II). Finalement nous examinons deux hypotheses sur la formation des bulles. The purpose of this work is to study some characteristics of roundish features which are observed on Hα filtergrams, above the solar chromosphere, at the limb. The ‘Bulles’ - chromospheric bubbles - so called for their apparent shape, move upwards into the low corona and fade out when their material is completely ionised (Figures 1 to 3). The physical data for nine observed bubbles listed in Table I are: lifetime (2), average height (3), ascending velocity (4), lateral velocity (5), angle of the movement with the normal to the solar limb (6), intensity in Hα (7) and diameters (8). The time variations of some of these properties are pictured in Figure 4. It is noted that all measured velocities are super-sonic. We discuss the relation between the ascending velocity and the maximum height reached by the bubbles (Figure 5) and compare it with the one obtained by Rush and Roberts for spicules. The study of their diameters shows that the vertical diameter is generally greater than the horizontal one. Their ratio is a function of speed and becomes equal to unity when the bubble velocity reaches 70 km s−1 (Figure 6). The bubbles seem to diffuse in the corona as spicules do. We also compare the properties of bubbles with some already known chromospheric features (Table II) and there seems to be a link between bubbles and Becker's grains. Finally, we, however, propose two hypothesis regarding the origin of bubbles: (i) they are related to spicules, or (ii) they are accelerated in the low chromosphere by the divergence of the magnetic field.

Spectral investigation of the chromosphere

Highly resolved Hα spectra and filtergrams obtained at the Fraunhofer Observatory on Capri were analysed by a method whose principles have been described before. As a result the tentative conclusion of our previous work has been confirmed: The mottles of the chromospheric fine structure are clouds superimposed on the low chromosphere. Furthermore, it is proposed that the latter is identical with the interior of the supergranular cells whose spatial averages lend themselves to an interpretation in terms of a spherically symmetric model. On the other hand, we present evidence that the boundary regions which are formed by the mottles of the fine structure pattern and which constitute the upper chromosphere are not adequately described by spatial averages. Instead, the properties of the individual structural elements should be taken into account by an appropriate theory, if only in a statistical fashion. Some of these properties, viz. source function, optical thickness, Doppler width and velocity, were measured for a large number of elements. On the basis of indirect evidence we suggest that the spicules seen on limb photographs of chromospheric lines are identical with the elements of the supergranule boundary structure.

EUV emission, filament activation and magnetic fields in a slow-rise flare

The evolution of coronal and chromospheric structures is examined together with magnetograms for the 1B flare of January 19, 1972. Soft X-ray and EUV studies are based on the OSO-7 data. The Hα filtergrams and magnetograms came from the Sacramento Peak Observatory. Theoretical force-free magnetic field configurations are compared with structures seen in the soft X-ray, EUV and Hα images. Until the flare, two prominent spots were connected by a continuous dark filament and their overlying coronal structure underwent an expansion at the sunspot separation rate of 0.1 km s−1. On January 19, the flare occurred as new magnetic fields emerged at ∼ 1019 Mx h−1 beneath the filament, which untwisted and erupted as the flare began. The pre-flare coronal emissions remained unchanged during the flare except for the temporary addition of a localized enhancement that started 5 min after flare onset. EUV lines normally emitted in the upper transition region displayed a sudden enhancement coinciding in time and location with a bright Hα point, which is believed to be near the flare ‘trigger’ or onset point. The EUV flash and the initial Hα brightening, both of which occurred near the center of the activated filament, were followed by a second EUV enhancement at the end of the filament. The complete disruption of the filament was accompanied by a third EUV enhancement and a rapid rise in the soft X-ray emission spatially coincident with the disappearing filament. From the change of magnetic field inferred from Hα filtergrams and from force-free field calculations, the energy available for the flare is estimated at approximately 1031 erg. Apparently, changes in the overlying coronal magnetic field were not required to provide the flare energy. Rather, it is suggested that the flare actually started in the twisted filament where it was compressed by emerging fields. Clearly, the flare started below the corona, and it appears that it derived its energy from the magnetic fields in or near the filament.

Umbral oscillations and penumbral waves in H?

We present examples of umbral oscillations observed on Big Bear Hα 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.

The fine structure of the solar atmosphere in the far ultraviolet

High resolution spectroheliograms in the ultraviolet emission lines He i, He ii, O iv, O v, and Ne vii have been photographed during a sounding rocket flight. Simultaneously, broad band filtergrams of the far ultraviolet solar corona were obtained from the same flight. This paper describes qualitatively the spatial distribution of the UV emission. A comparison with an Hα filtergram is made. The most significant results can be summarized as follows: We find most of the ultraviolet emission concentrated around spicules, with different degree of concentration, decreasing with higher temperatures. 4 different areas of ultraviolet emission can be distinguished. (1) The normal network, bright in all UV emission lines from the chromosphere into the corona. (2) The coronal holes, bright in all UV emission lines up to 600 000 K but depressed in coronal lines from 1 million degrees upward. (3) The coronal depressions near active centers, absence of all ultraviolet emissions and (4) Active regions, where ultraviolet emission comes from plages, sunspots and coronal loops. High non-thermal Doppler velocities can be found in certain plage kernels around 105 to 2 × 105 K. Sunspots are bright in the ultraviolet, but do not exhibit He i or He ii emission. The corona above sunspots is weak. Sunspots do not show high non-thermal Doppler velocities. The He i and He ii emission does not follow either chromospheric, transition zone or coronal pattern; one can recognize some typical behavior of each.

Magnetic field and electric current structure in the chromosphere

Three dimensional vector magnetic field structure throughout the chromosphere above an active region is deduced by combining high resolution Hα filtergrams with a simultaneous digital magnetogram. An analog model of the field is made with 400 metal wires representing fieldlines which are assumed to outline the Hα structure. The height extent of the field is determined from vertical field gradient observations around sunspots, from observed fibril heights and from an assumption that the sources of the field should be largely local. After digitization the magnetic field H matrix is retrieved. Electric current densities j are computed from j=curl H. The currents (typically 10 mA m−2) flow in patterns not similar to observed features and not parallel to magnetic fields. Lorentz forces are computed from {ie0323-01}. The force structures correspond to observed solar features and a series of observed dynamics may be expected: downward motion in bipolar areas in lower chromosphere, an outflow of the outer chromosphere into the corona with radially outward flow above bipolar plage regions (where coronal streamers are observed) and motions of arch filament systems. Observed current structure and magnitude agree well with previous vector magnetograph observations but disagree with theoretical current-free or force-free concepts. A dynamic chromosphere with electromagnetic forces in action is thus inferred from observations.

The galloping chromosphere

Oscillating velocity fields can be observed on Hα filtergrams as a shifting pattern of intensity fluctuations known as ‘the galloping chromosphere’. The characteristics of this activity are those of horizontal running waves of typical period (∼ 300 s) and long wavelength (∼ 20000 km) that can be interpreted as acoustic-gravity waves propagating in the acoustic domain. Periods are longer in dark, structured regions, and in fibrils, and the change is quantitatively consistent with the reduction of resonance frequency in a magnetic field of 1 to 10 G. These easily observed fluctuations thus offer a means of estimating magnetic-field strength at specific locations in the chromosphere.

Chromospheric Oscillations in Plages

The 5-min oscillations in Hα plages as reported by Bhatnagar and Tanaka (1971, Solar Phys.24, 87) are discussed with respect to spatial – both horizontal and vertical – phase relation. Data for analyses include (a) Hα filtergrams at −0.5 Å (1050 min in total) (b) Hα filtergrams simultaneously taken at ±0.5 Å (60 min) (c) Simultaneous Hα−0.5 Å and FeI 5233 Å videomagnetograms (Doppler mode) (100 min in total). Projection (16 and 35 mm) -photocell method (Bhatnagar and Tanaka, 1971) was used for tracing the data with special attention to the guiding and reproducibility. Also photographic prints were used for comparison. The results are as follows: (1)those showing oscillatory brightness changes are dark dot-like features (size ≲ 1″) in brighter regions of plage and slightly elongated features (1 × 1.5–2″) in less bright regions. The mean number density of these features is one per 2200 ± 500 km square.(2)There are always a few regions with a size 8″ or more in one plage which show very regular oscillations with a period of about 5 min for a long time (at least 6.5 h, Figure 1). The distribution of the intervals between successive intensity peaks is much narrower than the Gaussian distribution and has a skewness towards longer period. The other places in the plage show relatively irregular patterns, but with the mean interval of the fluctuations still being 5 min.(3)In regular 5-min brightness oscillations the phase of oscillations in one wing is complementary to that in the opposite wing, indicating velocity oscillations. When the intensity records are relatively irregular there are many cases (76%) in which the blue wing patterns can be matched well for the whole time sequence with the red wing records by advancing the red wing records by 196 ± 18 s, (see Figure 2) indicating that the observable dark feature first appears in rising mode and then with a constant time lag (close to the free fall time from −0.5 Å to +0.5 Å) appears in the falling mode.(4)The records of brightness variations are very similar when the aperture is changed from 4″ (minimum reliable aperture) to 10″ in most places. Comparison of the oscillation phases in all the measured points in one plage shows that a plage can be divided into several areas (for a certain time) with sizes from 7000 km to 30000 km in each of which the phases are very close to each other for 2 to 10 cycles. Figure 3 shows the distribution of the intensity peaks of oscillations measured at Hα–0.5 Å (the abscissa is time, the ordinate indicates space points arranged geometrically). The boundaries of these areas vary after the coherent cycles. The iso-intensity curves of the intensity-time records (see Figure 3) have slopes corresponding to horizontal phase velocity 100 km-600 km s-1.(5)The phase relation between Hα regular 5 min and FeI 5233 Å 5-min velocity oscillations at the same place is almost constant for more than 5 cycles although the phase lag is different from place to place. (See, for example, the period from 20:00 UT to 21:00 UT in Figure 1. The smoothed-out curve denotes the record of the velocity oscillations in λ 5233 Å taken simultaneously.) The distributin of the phase lags is peaked at about 1 min. Assuming 1000 km for height difference of two lines the vertical phase velocity is equal to 17 km s-1.(6)There is an indication that the flare originating in the plage occurs about 20 min after the oscillations start to be in phase (1-min accuracy) in very large scale (≳20000 km) although the samples (6 flares in three different plages) are not large enough for statistically meaningful discussions.

Evolution of Structures in the Bright Hα Network

We present off-band Hα filtergrams showing the bright Hα network (solar filigree) and its evolution, with subarcsecond resolution. In quiet regions the total area covered by elements in the network is constant over periods of hours but changes in the geometry are observed over 10 min (the timescale of granule and dark mottle lifetimes), that can be interpreted as due to (macro)turbulent horizontal velocities of about 1 km s-1. The network can be resolved into strings of tiny bright dots; on the order of 100 per supergranule cell boundary. These dots may mark the location of magnetic filaments. The contrast of the features is higher in the blue than in the red wing which may indicate that in the lower chromosphere, the downdraft at supergranule boundaries does not occur in the bright crinkles themselves but rather adjacent to them.

Visual Determination of the Mean Effective Chromospheric Thickness

IN the course of taking submillimetre chromospheric data from Concorde 001 during the eclipse of June 30 (ref. 1), it proved possible to record information on the visual thickness of the chromosphere which is accurate enough to merit reporting. During an eclipse viewed conventionally from a stationary site on the ground, the precise time at which the solar photosphere reappears after totality, termed third contact, may be predicted to within 0.5 s (of time), given an adequate knowledge of the lunar limb topography, that is, to within 1 km in height at the appropriate point on the limb. Further, photometric measurements in white light could determine this time to within 0.3 s. It may also be possible from the ground, under ideal observing conditions, to see and measure emission from the chromosphere which should be visible in isolation for around 3 s, just prior to third contact. The predominant radiation seen during this time is that of the first line of the Balmer series, Hα, at 6,563 A which accounts for the traditional description of the phenomenon as the ‘red flash’. It should be possible, after taking into account atmospheric scattering and diffraction effects, to measure the duration of this to within 0.5 s, though measurements of this type are not typically considered sufficiently accurate to be worth taking, when, for example, the flash spectrum of the chromosphere is also available2. Although various observers3,4 have studied the chromosphere spectrally in emission during eclipse, both just after second contact and just before third contact, inaccuracies in our knowledge of the precise orientation of the Moon and its topography have made absolute heights above the photosphere very difficult to measure to better than 300 km, although relative heights may be estimated to within 100 km under ideal seeing conditions. Similar restrictions due to seeing apply to estimates based on Hα filtergrams taken outside eclipse conditions5.

Lifetimes of enhanced chromospheric network features near active regions

Centerline Hα filtergrams providing nearly full day coverage of the Sun are used to study the lifetimes of enhanced network features near active regions. In the two cases studied the fraction remaining of those features present at an original epoch remains near unity for 50 h, then drops exponentially with a 1/e decay time of 30 h. Histories of representative enhanced network features are discussed.

The dynamics of solar surges

Curves of velocity as a function of height have been determined from time sequences of on- and off-band Hα filtergrams of four surges. The surges behave as predicted by the melon-seed mechanism. However, in the outward-moving material the deceleration taking place after the maximum velocity has been reached, is stronger than if gravity alone is acting in three of four cases; during the return phase, the acceleration toward the solar surface is less than free-fall in two of three cases. The braking force, opposing the motion, is proportional to the velocity of the surge material. Non-force-free effects in the magnetic field supporting the surge provide a reasonable explanation for the braking.

Some statistical properties of Ellerman bombs

One hundred seventy-eight Ellerman bombs were identified and studied with on- and off-band Hα filtergrams of two active centers, one near disk center and the other near the east limb. The photographs, taken through a 1/4 A bandpass filter, occasionally attained resolution as fine as 0.3″. The mean duration of bombs at Hα-2 A was about 13 min near disk center and 11 min near the limb; these times increase slightly when we observe closer to the core of Hα. Eighty-six percent of the bombs in the near-limb region and 56% in the disk-center region were seen to be accompanied by ejections of dark material; the ejections were 3–30″ long. The ejection length appears to be proportional to the bomb size times the bomb lifetime.

The topological association of H? structures and magnetic fields

The topological associations between Hα structures and magnetic fields are examined for an active region observed on two different dates. The structures seen in the on and off band of Hα filtergrams are compared with the contour maps of magnetic fields at the level of magnetogram observations. Similar comparisons are made also with the configurations of force-free magnetic lines of force at various heights evaluated with the use of formulations developed previously by Nakagawa and Raadu (1972).

On some transient H? features associated with metric type III bursts

The question of the chromospheric features type III association is reconsidered by using Hα filtergrams both on and off band. A set of 44 metric type III groups, for which an association can be ascertained with a high degree of confidence leads to the following results:

Observations of moving magnetic features near sunspots

The properties of small (< 2″) moving magnetic features near certain sunspots are studied with several time series of longitudinal magnetograms and Hα filtergrams. We find that the moving magnetic features: (1) Are associated only with decaying sunspots surrounded entirely or in part by a zone without a permanent vertical magnetic field. (2) Appear first at or slightly beyond the outer edge of the parent sunspot regardless of the presence or absence of a penumbra. (3) Move approximately radially outward from sunspots at about 1 km s−1 until they vanish or reach the network. (4) Appear with both magnetic polarities from sunspots of single polarities but appear with a net flux of the same sign as the parent sunspot. (5) Transport net flux away from the parent sunspots at the same rates as the flux decay of the sunspots. (6) Tend to appear in opposite polarity pairs. (7) Appear to carry a total flux away from sunspots several times larger than the total flux of the sunspots. (8) Produce only a very faint emmission in the core of Hα. A model to help understand the observations is proposed.