H-alpha observationprogram



Ever since Coronado has been marketing high-quality and relatively inexpensive H-alpha filters and telescopes, more and more amateur solar observers dispose of their own H-alpha equipment. This trend is also noticeable within the Belgian Solar Section (Belgian Astronomical Association, VVS).

Hence the need for a H-alpha observation program. On the one side, it would be good that everybody disposing of a H-alpha instrument can step into this program in a simple manner. On the other side, the program needs to be flexible enough to allow for an expansion. Moreover, because the observations should be comparable with other international solar sections, it is desirable that the same observation criteria are being used. Finally, it would be nice if contributions can be made to the professional solar program.

From this, it seems appropriate to construct a program similar to that of white light observations. There, the basic consists in the determination of the Wolfnumber, and an advanced program exists for the determination of the hemispheric Wolfnumbers, the Becknumber, the Classification Value, positions of sunspotgroups and polar faculae numbers.

Translated to H-alpha, the basic program could be the counting of the prominence hearths. A more advanced program could contain the determination of the number of prominences per hemisphere or per zone (polar/main zone), the determination of the prominencenumber, the subdivision according to size, the classification according to Zirin, and latitude determination. Potentially, a professional contribution could consist in the determination of the optical classification and other features of solar flares, in support of X-ray observations.

It is obvious that, just as for white light observations, some observers can further specialize. Amongst other are the imaging of the aforementioned solar phenomena via (highresolution) pictures or even videotapes, and studies of chromospheric faculae and filaments. For the moment, these topics will not be discussed here.

Paragraphs underneath will deal with the observation of the chromospheric phenomena and related criteria.

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Image quality

Before starting the observation, it is good to get an idea of the current atmospheric conditions. Quantification is done by using the seeing parameter Q. This is a measure for the atmospheric instability and is also used for visual observations of the photosphere. The scale is also employed by the SIDC, in steps of 0,5. In practice, the value of Q as deduced from white light observations is also applicable on H-alpha observations. If no observations of the photosphere are carried out, the same approach can be used by looking at the sharpness of the chromospheric network and the spicules (chromospheric edge). See table underneath.

QDescription white lightDescription H-alpha
5Near-perfect image. The granulation is visible over the entire solar surface. There are no vibrations at the solar edge. The fine structure in the penumbra can be well observed.Near-perfect image. The details in the chromospheric network are well visible over the entire solar disk. There are no vibrations at the solar edge. Even the finest spicules are visible at the solar edge. The finestructure in filaments is well observable.
4Very small groups and faculae are visible. The granulation is well visible.Very small details in the filaments and the chromospheric faculae are visible over the entire solar disk. The chromospheric network is well visible. The solar limb is moving slightly.
3Also the smaller spots are visible. The shape of the penumbrae is well visible. The largest faculae can be seen as well.Chromospheric network is visible over the entire solar disk. Spicules, just as the smaller filaments are visible. The shape of the chromospheric facuale is well visible. The solar limb is vibrating a bit, but this vibration is hardly noticeable in the umbrae.
2Only the biggest spots can be observed. Both umbrae as penumbrae are visible. Only the big filaments are observable. Hints of chromospheric network. The chromospheric faculae melt into the chromospheric background. The solar limb is moving strongly, and this movement is also visible in the umbrae of the sunspots.
1Only the biggest spots are visible. Penumbrae are not visible, only the umbrae.Only the biggest and darkest filaments can be observed. Umbrae of the sunspots are not or hardly not visible. The entire image is waving.

However, while observing prominces, also the transparency of the air should be taken into account. During periods of bad transparency (e.g. haze), the background in the eyepiece's field of view will turn red, making it very hard to observe small prominences or details in prominences. In 1968, B. Wedel [1] introduced a transparency scale W that need to be used while carrying out such observations. See table and figure underneath.

1Celestial background is very dark, the prominences are very distinct.
2Celestial background is dark, the prominences are distinct.
3Celestial background slightly bright, but the prominences are still quite distinct.
4Celestial background is bright, the prominences can only be seen with difficulty.

In practice, Q usually has the value 3 and W about 2,5. Observations with Q < 2 and W = 4 are to be avoided. Observation with Q = 5 and W = 1 are exceptionnal, unless one observes from high mountain tops. The moments with the best seeing/transparency vary with location, the seasons and of course the day-to-day weather conditions.

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Orientation of the solar image

The orientation of the solar image is done the same way as in visual observations of the photosphere ([2]). See attached figure.

  1. Put the sun as centrally as possible in the field of view. Choose an ocular allowing a full view of the sun.
  2. Let the scope immobile (turn of the motor drive when a motor-driven telescope is used). Because the sun has an apparent daily motion from the east to the west, she will not stay in the view's center, but will gradually move through the field of view.
  3. The line along which the sun has been moving, can now -in thoughts- be drawn over the center of the solar disk. If the position of visible sunspots is well known, it can be used as an additional reference. The leading solar edge (the edge which moved first out of view) is the western edge, the other is the eastern edge. Because of the rotations of sun and earth, as well as the movement of the earth around the sun, new filaments and prominences will first appear at the eastern solar limb, and will rotate off the solar disk about two weeks (half a solar rotation) later at the western solar limb.
  4. The east-westline divides the visible solar surface in 2 areas: the northern and southern hemisphere. To know which hemisphere is the southern, put the sun again as centrally in the field of view as possible. Move the (front of the) telescope slowly down. The part of the sun that stayed the longest in view, is the southern hemisphere. The part that moved out of view as first, is the northern hemisphere.
  5. By tracing a line perpendicular to the east-westline and through the centre of the solar disk, one defines the north-south-axis. The east-westline and the north-south-axis are always perpendicular to each other. However, the positions of these lines is only apparent, because...
  6. The axis of rotation of the sun stands, just like that of the earth, oblique to the surface described by the rotation of the earth around the sun. As a result, the real solar axis makes an angle with the apparent north-south-axis. Moreover, this angle varies throughout the year. This angle is called the P-angle, and varies between -27° and +27°. Theangle is counted positively from the north towards the east. Its value can be found in an astronomical almanac, on websites as BASS2000, or it can be calculated (see Calculation of solar paremeters for the necessary formulae). For the orientation of the solar image, an accuracy of 1° is amply sufficient. In the example, the correction was carried out for P = +20°.

Due to the rotation of the earth, the positions of the east-westline and the north-south-axis vary thoughout the day. This is why for each observation, these 2 lines need to be determined everytime again. The use of a crosswire ocular increases the accuracy considerably. At all times, care has to be taken north and south are not reversed, especially if another telescope is used (e.g. for observing sunspots) where north and south are mirrored.

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Observation programs

Prominence hearths
ProminencenumberSolar hemisphereZoneClassificationSizeLatitude
Solar flares

Prominence hearths

Prominences, just like sunspots, group together. Therefore, the simplest observation program consists in counting the number of prominencegroups. These groups are also called hearths (from the German "Herde").

Prominences are taken together as soon as they are 5° or less in latitude from each other. If the distance between prominences is bigger, they can still be taken together provided there seems to exist a physical link. These "rules" were deduced from 2 drawings made by the German solar section Sonne ([3] and [4]). See figure underneath for examples. The British Astronomical Association seems to use this parameter in the determination of prominence activity (PMDF = Prominence Mean Daily Frequency), and the German Sonne uses this parameter as a base for its Prominencenumber Rp.

In the figure, 6 prominence hearths are visible. The second prominence from the right is located 6° from the first, and 7° from the third prominence. Because it does not appear to have a physical link with its neighbours, it is considered as a separate prominence. In prominence four this link seems to be present. So, despite the fact that both ends are located 8° from each other, they are considered as 1 prominence. A more evident example can be found in prominence five, where the ends are located even further from each other, but the physical link is also much clearer (the parts of the prominence almost touch each other). In hearth six, the prominence distances are 4° and 3° respectively, hence they form one group. A prominence which is detached from the solar limb is considered a fullgrown prominence.

Figure underneath provides an idea about what a subdivision in steps of 10° looks like. 5° is a really small distance (a quarter of a circle circumference divided by 18). It is assumed one can perform an accurate subdivision based on visual means only, though a crosswire is of course a great help.

Counting prominencegroups can lead to interesting insights. From observations made by the British Astronomical Association (PMDF), it appears that maximum PMDF and its evolution during solar cycle maximum precedes the evolution of flare activity by 1 to 2 years. It may also be clear that during cycle minimum, there always are a few prominences visible. The rate cycle minimum / cycle maximum is approximately 1 for 30 for solar eruptions, 1 for 15 for the Wolfnumber R, and 1 for 3 for the PMDF!

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A possible expansion of the basic program consists in determining the prominencenumber Rp. It is developed similary to the Wolfnumber R. For the determination of R, the number of sunspotgroups "g" and the number of individual spots is counted, for Rp the prominence hearths "H" and the individual structures "e" (from the German "einzel") are counted. Also the formulas are comparable. For the Wolfnumber it is R = 10.g + f, and for the prominencenumber Rp = 10.H + e.

Though no restrictions in size or lifetime were found in the literature, prominences should not be confused with spicules. These are 500 km wide, threadlike, up and down moving structures existing for about 10 minutes and reaching heights between 3.000 and 9.000 km ([5], [6], [7], clip (Movies; clip 7)). They are caused by oscillations of the solar surface (Bart De Pontieu, [28]). Because of their height (1% of the solar radius), it seems appropriate not to count any prominences smaller than 1% of the solar radius. In practice, this means that a small structure is considered a prominence if it stands out clearly and at least for 15 minutes above the solar limb.

There does also not exist a restriction on the type of prominence. Prominences which are detached from the solar limb, are counted all the same. NO protrusions, twists, brightenings,... in a prominence are counted separately.

In the example above there were 6 prominence hearths, with per hearth respectively 3, 1, 4, 2, 4 en 3 (= 17 in total) individual structures. The prominencenumber is thus Rp = 6.10 + 17 = 77.

The prominencenumber Rp was developed by Völker in 1970. Strangely, the system seems only used by the German section Sonne itself. Graph underneath gives an idea of the evolution of Rp during the last two solar cycles ([8]).

© Stetter H., Sonne 105

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Per solar hemisphere

Because the solar activity (Wolfnumber, solar eruptions) evolves differently per solar hemisphere, it seems straightforward to make a similar distinction also for the prominence hearths and Rp.

The method to split the sun in a northern and southern hemisphere was already explained in the orientation of the solar image. The determination of the number of hearths and the prominencenumbers is done in the same way as described above. The numbers get an index "N" or "S" to indicate on which hemisphere they are applicable. In principle, it concerns HN, HS, RpN and RpS, for which of course H = HN + HS, and Rp = RpN + RpS.

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Per zone

In contrast to the sunspotgroups, which seldomly attain latitudes above 40°, prominences do appear at these higher latitudes. According to Sonne, the evolution of these polar prominences differs entirely from those from the main zone ([9], [10]):

Hence, a subdivision between polar and main zones seems necessary. Unfortunately, things get a bit more complicated here because there exist no unanimity concerning the border between the two zones. Sonne ([10]) uses 50° latitude as a border, while "The Astronomer" and BAA use 40°. The difference probably originates in the interpretation of "main zone". If one only considers sunspots, then the zone between +40° and -40° is to be considered as the main zone. This is because during the last 3 solar cycles only about 30 groups (of almost 11.000!) appeared at an higher latitude. On the other hand, sunspots do appear at latitudes above 40°, and faculae fields -the precursors of sunspotgroups- appear up to 55° (above that latitude polar faculae appear; [13]). In view of the mentioned physical connections, a subdivision at 50° seems reasonable. In case of doubt (e.g. in case of a prominence stretching from +45° to +55°), the prominence can be considered as belonging to the polar zone.

For the visual determination of these zones, firstly a split into solar hemispheres is being made as described above. To define the two zones visually, following ways exist:

Figures underneath show the subdivision per zone, and illustrate also the methods to visually define the the two zones.

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There exist 2 major classification schemes.

The first one was introduced by Völker in 1969. It makes a difference between the individual prominences based on the height (small, big or unusually big), the outlook (bar, arch or area) and if the structure is detached of the solar limb or not ([14]). The overview is given in figure underneath and originates from the website of Anthony & Christopher Seal.

© Anthony & Christopher Seal

The problem with this classification is that it needs to be applied on each individual structure within a prominencegroup. This is very timeconsuming, and one could question also the benefit (except for the height). It is also significant that nowhere reports can be found on this classification, not even in Sonne.

The second scheme is based on the classifiction by Zirin ([15]). Here, a distinction is being made between "quiescent" and "active" prominences, respectively called QRF en ARF. QRF can be subdivided in 5 types (A to E), and ARF in 4 types (F to I). It are especially the ARF and type E that matter, because here material can be ejected into space. In particular, when these particle clouds are ejected on the westside of the sun, they can have a significant geomagnetic effect. Figure underneath comes from the website of David Knisely.

© David Knisely

This classification has the advantage it only needs to be applied on the entire prominencegroup. Moreover, there is a real physical link with chromospheric activity. Unfortunately, also for this classification no reports are available.

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Per size

Prominences can also be subdivided according to their size. This can be done in 3 ways: according to the height reached, according to their area, and according to the length of the corresponding filaments.

  1. Height
  2. In fact, this concerns the size-determination according to the system developed by Völker. Regardless the type, the prominence hearth is subdivided in "small" ("s"), "large" ("l") and unusually large ("xl"). This corresponds to heights above the solar limb of respectively (expressed in the radius of the sun Rsun):

    Figure underneath provides an idea on how high these prominences need to be in order to belong to one of the three mentioned categories.

    According to Zirin ([16]), there is an eruption of every prominence (within 48 hours) if its height exceeds 50.000 km, this is 1,07 . Rzon.

  3. Area
  4. Professional astronomers (Catania) use the "Prominence Profile Area" (PPA) to track the prominence activity ([17]). On a H-alpha picture, the solar limb is subdivided in 360 areas of 1°. The number of arcseconds above each area of 1° wherein a prominence is present, is counted. Finally, all areas are taken together and summed. Of course, this is a very timeconsuming occupation.

    An alternative was developed by Völker ([18]). The method is similar to the area-determination for faculae or sunspotgroups. A transparent with mm-paper is put over the picture, and the number of mm2 that is completely filled with prominences is counted. The sum gives the area A. One has to pay attention that this determination is always done for the same solar diameter. Because this one changes in the course of a year (movement of the earth around the sun), a correction factor needs to be applied. The accuracy of the method (and also the work charge) increases with the size of the solar image, which needs to have a minimumdiameter of 15 cm.

    A final method is applied by "The Astronomer" ([19]) and consists in subdividing the solar limb in areas of 5°. Each area in which a prominence is present, gets the value "1". Finally, all areas are summed. Not all prominences are eligible for the count. E.g. if an area (of 5°) is skipped by a large, archlike prominence, then this area is not counted. A prominence thus needs to be attached to the solar limb, and prominences should not belong to Völker's type "detached". Also short-lived, eruptive prominences should not be counted. On pure visual grounds, this method does not seem accurately applicable. Based on a picture (or using a detailed crosswire ocular), the counting can be quickly executed.

    A study ([20]) of quiescent prominences (QRF) has revealed there exists an excellent correlation between the prominencenumber Rp as determined by Sonne, and the Prominence Profile Area (PPA) as determined by Catania: 96%!

  5. Length of filaments
  6. According to [21], filaments having a length of more than 15° have a higher chance to produce a geomagnetic impact (if they explode) than smaller filaments. Observing prominences can not rpovide a clearcut answer on their true length, as may be clear from sketch underneath.

    There exists several programs measuring the length of quiescent filaments (QF; not related to sunspots), but they are all timeconsuming and actually require the use of pictures ([22]). An alternative may be to estimate the length of each QF and subdivide in bigger or smaller than 15° length. While estimating their length, one needs to take into account the foreshortening: filaments at the solar limb are longer than they appear (just like the length of sunspotgroups). Figure underneath can be useful: It shows circles met radii of respectively 15°, 30°, 45°, 60° and 75°.

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Latitude determination

The latitude of each prominence hearth can be determined by using a crosswire ocular or more accurate with a picture. Some possibilities are discussed in the Solar Astronomy Handbook ([23]). Though all determinations are time-intensive, the results do give a detailed idea on the evolution in number and latitude of the prominences. Figure underneath gives the procentual number of prominences as a function of the latitude, for the years 1999 to 2002. The physical events at the solar poles (see Per zone) become more clear here.

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Solar flares

The observation program for solar flares is mostly a copu paste of the program developed by the Australian Spaceweather Bureau (IPS; [24]), supplemented by some of my own comments. A similar program is described in [25]

Solar flares are transient brightenings of the chromosphere that may last for a few minutes to a few hours. The energy released in a solar flare may reach up to 1025 joule (about 10 billion atomic bombs of 250 kTon each!). Most flares are only visible by observing various spectral lines with a narrow band filter. However, the highest energy and brightest flares may by seen in unfiltered white light. These "white light flares" (WLF) however, are very rare, and can be expected to occur on average about 5 times per solar cycle. A broad band filter in the blue part of the spectrum (such as a Kodak Wratten #47) will substantially improve the chances of seeing a solar flare in continuum light, and at times of solar maximum, such a filter may allow the observation of a dozen flares or more per year.

Typically the hydrogen-alpha spectral line, at a wavelength of 656,3 nanometres is used for solar flare patrol. A filter with a passband of 1/4 to 1/2 Ångstrom (0,025 - 0,05 nm) is commonly used for this application. A tunable filter is an asset to examine chromospheric motions associated with many flares, and the spectral width of a flare may also be used to indicate the brightness category of a flare if accurate photometry is not available.

In general flares will occur close to an active region (sunspotgroup), but occasionally a flare may occur well away from an active region, but associated with a large dark filament. These are termed Hyder flares.

There are two aspects to flare patrol: detection of a flare, followed by characterisation of the flare. Although a flare is defined as a transient brightening of the chromosphere, there are forms of chromospheric brightening which are not classed as flares. These are related to both the brightness and the area of the brightening, and it is these two characteristics which define when a flare has occurred.

To be classified as a flare, the chromospheric brightening must exceed an area threshold and a brightness threshold. For H-alpha observations, the brightness threshold is 150% of the chromospheric background, and the area threshold is ten millionths of the visible solar hemisphere. A brightening which has less than 150% of chromospheric background (even though it may have more than 10 millionths area) is termed a plage fluctuation (a plage is a chromospheric faculae field). A small brightening which exceeds 150% of chromospheric background but less than 10 millionths hemispheric area is termed a point brightening.

Once a flare has been detected it must then be characterised:

ObservationChromospheric brightening
Detection limitsArea > 10MHBrightness > 150%
Primary characteristicsBeginArea
Secondary characteristicsActive region or Hyder flare?Dark surge?
Coverage Umbra?Disappearing filament?
Parallel (//) of confluent (Y)?Doppler shift?
Loop system?Sympathetic (simultaneous) flares?
Single- or multiple eruption centers?Moreton wave?
Number of brilliant points?Visible in white light?

There are six essential characteristics by which a flare is usually described. To these may be added several associated characteristics. These secondary features are usually added because of the information they potentially convey about possible geoeffective events that might be associated with the flare. Some of these may point toward associated coronal mass ejections; some may indicate high temperature flares which typically are associated with high energy particle emission (protons,...).

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Primary flare characteristics

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Secondary flare characteristics

There are many secondary characteristics that may be associated with a flare and different organisations choose to report varying subsets of these. Some parameters indicate high temperature flares, some indicate associated coronal mass ejections, and others have questionable correlations with any significant effects. Some of the potentially more useful are discussed below.

A large parallel ribbon flare - Learmonth Solar Observatory

Basic document by John Kennewell. © Copyright IPS - Radio and Space Services.

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Observation report

Reports of H-alpha observations need to satisfy the same requirements as the visual observations. The (monthly) report needs to be send as quickly as possible after the end of each month to the section leader (Franky Dubois), and has to contain essential elements like observing equipment , the address of the observer, and a short explanation of the symbols used (e.g. gN is the number of sunspotgroups on the northern solar hemisphere). It is best to make it immediately in English, so that the same report can be send to other international solar sections.

For H-alpha observations, the aim is to transmit as much data as possible through electronic mail (e-mail) and in table-format (excel-file). This facilitates the processing work considerably. Any comments can be added in the cel itself ("comments") or underneath the table.

Though it is desirable that everybody uses more or less the same format, this is possibly not feasible because everybody has his own program. This consists of the basic activity supplemented by observation programs of personal preference.

Examples underneath give an idea how an observation report can be constructed. As can be seen, observing in H-alpha can lead to an entirely seperate observation discipline requiring a separate observation sheet. It is up to the observer to make this distinction.

Example 1: Basic H-alpha program supplemented to photospheric observations

Solar indices for the month of August 2005
Jan Janssens, Freesiadreef 3 B43, 1030 Schaarbeek BELGIUM e-mail: j.janssens@chello.be
EquipmentCelestron8, 203.2 mm, f/10, 68*, 1000 Oaks Obj. filter, Blue OcfNaked EyeCoronado PST, < 1 Å, 40 mm, f/10, 33*
DayTimeQgngsCVnCVsPFnPFsAWProm NProm SRemarks

Example 2: Advanced H-alpha program

Solar indices for the month of August 2005
Jan Janssens, Freesiadreef 3 B43, 1030 Schaarbeek BELGIUM e-mail: j.janssens@chello.be
EquipmentCoronado PST, < 1 Å, 40 mm, f/10, 33*

Example 3: Advanced H-alpha program

Solar indices for the month of August 2005
Jan Janssens, Freesiadreef 3 B43, 1030 Schaarbeek BELGIUM e-mail: j.janssens@chello.be
EquipmentCoronado PST, < 1 Å, 40 mm, f/10, 33*

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  1. Beck et al., Solar Astronomy Handbook §B.9.4.1, Willmann-Bell Inc., 1995
  2. Taylor P.O., Observing the Sun pp. 73-74, Cambridge University Press, 1991
  3. Beck et al., Solar Astronomy Handbook Fig. B.9.17-18 pp. 344-345, Willmann-Bell Inc., 1995
  4. Roth G.D., Compendium of Practical Astronomy Vol. 2, Fig. 13.23-24 pp. 36-38, Springer Verlag, 1993
  5. Zirin H., Astrophysics of the Sun, pp. 161, 169, 187, Cambridge University Press, 1988
  6. Solar Terrestrial Dispatch, Internet Spaceweather Forecasting Course §2.3, 1996-1999
  7. Beck et al., Solar Astronomy Handbook pp. 323-324, Willmann-Bell Inc., 1995
  8. Stetter H., Sonne 105, pp. 8, VdS, Maart 2003
  9. Roth G.D., Compendium of Practical Astronomy Vol. 2, pp. 39, Springer Verlag, 1993
  10. Stetter H., Sonne 101, pp. 10-11, VdS, Maart 2002
  11. Altrock D., "Rush-to-the-Poles", NOAO, Maart 1999
  12. Phillips T., The sun does a flip, Science@NASA, Februari 2001
  13. Beck et al., Solar Astronomy Handbook §B.6.3.3, Willmann-Bell Inc., 1995
  14. Beck et al., Solar Astronomy Handbook §B.9.4.3, Willmann-Bell Inc., 1995
  15. Zirin H., Astrophysics of the Sun, Table 9.1 pp. 266, Cambridge University Press, 1988
  16. Zirin H., Astrophysics of the Sun, pp. 267, Cambridge University Press, 1988
  17. Beck et al., Solar Astronomy Handbook §B., Willmann-Bell Inc., 1995
  18. Roth G.D., Compendium of Practical Astronomy Vol. 2, §13.6.5, Springer Verlag, 1993
  19. Meadows P., Persoonlijke e-mail berichten, 28 juli 2005
  20. Stetter H., Sonne 72, pp. 236-238, VdS, December 1994
  21. Solar Terrestrial Dispatch, Internet Spaceweather Forecasting Course §12.5, 1996-1999
  22. Beck et al., Solar Astronomy Handbook §B., Willmann-Bell Inc., 1995
  23. Beck et al., Solar Astronomy Handbook §B.9.4.7, Willmann-Bell Inc., 1995
  24. Kennewell J., Solar Flare Patrol, IPS Radio and Space Services, ©2001
  25. Beck et al., Solar Astronomy Handbook §B.9.5, Willmann-Bell Inc., 1995
  26. Nitta N., Moreton waves and Coronal waves, Yohkoh Science Nuggets MSU, Februari 2002
  27. Zirin H., Astrophysics of the Sun, Fig. 11.8, pp. 390-391, Cambridge University Press, 1988
  28. De Pontieu B. and Nelson B., Scientists explain mysterious plasma jets on the sun, Lockheed Martin Solar and Astrophysics Lab, 29 July 2004

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Appendix 1: Crosswire ocular

The division in solar hemisphere or in zones is not always evident based on purely visual observations. Help by using a crosswire(ocular) is very welcome. Discussion underneath applies in principle only to the Coronado PST (40 mm aperture, 400 mm focal distance, 12 mm Kelner-ocular), unless mentioned otherwise.

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Explicit guidelines Belgian Solar Section

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Appendix 3: Comparison with BAA

E-mail of 5 November 2006 to BSO's solar section leader Franky Dubois. It was also published in the November 2006 Newsletter of the Solar Section.

After the solar section meeting early October 2006, I have had some correspondance with Marlyn Smith (BAA) on this topic. Just as a reminder: While I was developping the H-alpha program last year, I did contact BAA (Mike Beales) but without any response. At that time, I simply used what I could find on their website.

The appendix describes the way used by the BAA-members to determine the prominencenumber (their "MDPF" = our "H"). The most noticeable differences are:

I decided to check my observations (Dec 05-Mar 06, Jun-Aug 06) for any influence from the last 2 differences (detached, 3%). The eruptive prominences during this period were of no significance, I counted 1 at most. Of course, the influence of these eruptive prominences can become bigger during solar maximum. I do wonder why one would omit such prominences from the prominencenumber H: why count a "quiet" prominence in the morning, and then not count it a few hours later when it is erupting? After all, the prominencenumber is a measure for chromospheric activity, isn't it?

My initial average over the considered 7 months was 6,6, against 3,9 for the BAA. After my evaluation, my H-value dropped to 4,6: 1,4 was on the account of small prominences with heights between 1 and 3% of the solar radius, 0,6 came from the detached prominences.

Excellent news is that the overall evolution of H did not change after the evaluation: same evolution, but on the average 2 units lower. As far as I'm concerned, this is an indication not to change anything to the way we are currently counting prominences.

4,6 is still a bit higher than 3,9. In my opinion, this is due to the fact that I observe only a few days per month, and thus there could exist a certain bias. Also, I do see some prominences of medium height, but very lightweak.

Based on the above, I don't think we have to change something in the way we are counting prominences. Our monthly H-values will always be higher than those of the BAA, but now we know why this is, and also that we are in perfect agreement with the way Sonne is determining its prominencenumber.

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