The photosphere or the 'surface' of the Sun is the region where the temperature and pressure conditions of the plasma (hot ionised gases) first allow light to escape from the interior (become transparent to visible light) below that region it is opaque to light. The emission from the photosphere is incandescent meaning it is a continuous spectrum. However, this continuous spectrum is interrupted by many Fraunhofer absorption lines (figure 1).
Figure 1 The continuous 'black body' spectrum of the Sun containing the Fraunhofer absorption lines.
These dark lines (missing parts of the spectrum) occur because the Sun is composed of many elements including hydrogen, helium, calcium, iron, sodium and magnesium to name but a few. Due to the high temperature and pressure, these elements are in the form of a hot ionised gas. These elements absorb energy from the black body radiation at a very specific wavelength of light unique to them. This occurs when an electron jumps from an orbit with low energy to an orbit with a higher energy; the energy is absorbed from the spectrum and this produces the dark (absorption) line.
Above the photosphere is the chromosphere at around the height of 2-3,000km and the corona 10-20,000km, both are transparent to visible light; however the spectrum in these regions are dominated by emission lines. This is when the ionised gases reach a certain temperature and pressure where an electron from a higher energy state moves to a lower energy state and emits a photon. In the chromosphere, Hydrogen alpha emits this photon at a wavelength of 656.3nm, and Calcium II K (393.3nm) and H (396.8nm). Calcium II K and H however can be observed at a range of heights depending on whether the K1, K2 or K3 regions are observed (see figure 3&4). These electron shifts occur at a very specific temperature and pressure which is unique to each element, in turn this is directly linked to the height above the photosphere in which these conditions are right. We observe these emissions from a height determined by the opacity. As Calcium II K and H and Hydrogen alpha create different degrees of opacity, we see their emissions (and the structures revealed by their illumination) at different heights in the chromosphere.
The chromosphere is by no means a homogenous body of gas either. Spicules are huge columns of plasma rising and falling every 15 minutes and are present up to a height of 9,000km; these are drawing energy upwards and contributing to the heating of the chromosphere and corona. The idea of fixed temperature and pressure zones are probably not entirely the case. The basic knowledge of the structure of the chromosphere and corona and how it is heated is mostly unknown (figure 2).
Figure 2 The temperature increases in height above the photosphere.
So, after my brief summary, what is really going on with the physics? I asked a professional solar astronomer Dr Matt Penn (National Solar Observatory, Tucson, USA) a few questions for us.
Why do we see a surface when we look in hydrogen alpha and Calcium K?
“One way to think of chromospheric imaging is to approach it from the perspective of the ‘surface of last scattering’. We see the photosphere at continuum wavelengths when a photon leaves the surface of the Sun and travels to Earth. We also see the chromosphere when we look at a photon at H-alpha wavelength when it finally breaks through the haze and travels to Earth too.
The opacity of the gas at the Sun determines where these last scattering surfaces are located, and the opacity is highly dependent on wavelength. In the continuum the gas opacity is determined by H- bound-free atomic physics (interactions of two electrons with an H atom), but in the spectral lines it's bound-bound transitions of electrons, like you describe in your write-up. Electrons move from one level to another in a neutral atom as they get excited or relax.
So the relative strengths of these different types of opacities vary, depending on the physics of these processes. The bound-bound transitions have strong opacities. So at H-alpha wavelength, the photons leave the photosphere, almost to be immediately scattered by an H atom. It absorbs the photon, and then re-emits the photon in a different direction. The photon makes a random walk through the chromospheric gas, getting absorbed and re-emitted, until by chance it gets to a certain height where it is re-emitted towards the Earth, and does not encounter another H-atom. This height is the surface of last scattering, and defines what we see when we observe the H-alpha photons coming from the Sun.”
What heights in the solar atmosphere do these ‘surfaces of last scattering’ occur?
“Two main points are: the abundance of the particular atom in the Sun, and the opacity (or line strength) of the particular electron transition of that atom.”
Why do spectral lines have the shape that they have?
“It's the Doppler shift combined with the temperature of the gas. For a given temperature of gas, you'll have a certain fraction of the atoms moving at a certain speed, both towards and away from you, but there will be no bulk motion, so the mean speed will be zero. So if a photon is emitted from the photosphere at a wavelength blue of H-alpha, say H-alpha+10km/sec, on it's journey through the chromosphere it will encounter some H-atoms moving at that speed and so it can be absorbed and re-emitted by those atoms (and only those atoms). But there are fewer H-atoms at that speed than there are at zero velocity, so fewer of those photons will be scattered, and this is why the Sun looks brighter as you tune from line center to the continuum. Also, effectively you are moving downward towards the photosphere as you tune from line center to the continuum... in both red and blue wings. So the shape of the absorption lines depend on the distribution of speeds of the atoms (and their ability to absorb light because of the Doppler shift in their frame) which in turn is caused by the temperature of the Sun at the height that the bulk of the line is formed. This is why photospheric lines are narrow (they represent the photospheric temperature of 5880K) and chromospheric lines are broad (they represent the temp around 8000K). ”
Why do some absorption lines have emission cores in them?
“In the case of flares or of active region plages (in CaK) there is another form of energy, non-thermal wave energy or fast particles from a flare perhaps, which causes the electrons to get excited and then emit light at these wavelengths as they relax to the ground state. In three words: it's very complicated!”
Do the lines of Calcium K and H show the same features?
“I think the line profiles are nearly identical. Because the lines are so broad, there are several other absorption lines which appear on top of each of the H and K lines, and these lines are different between H and K of course. Also, the details of the polarization of the H and the K lines are different since they originate from different electron configurations. However in terms of the intensity spectra that most people would study, they would be the same.
So how do we observe these emission lines? To observe the chromosphere we can view it during a total solar eclipse when the visible spectrum from the photosphere is obscured by the moon and the bright emission spectra from the chromosphere can be seen. Another way is by using specific filters which can only pass the specific wavelength of light emitted by the element of choice. This happens also to be directly where the dark absorption line is in the photosphere. Therefore the photosphere is excluded and the faint emission spectra from the chromosphere can be observed instead.
Table 1 contains the most common elements found in the Sun and the point in the spectrum where the absorption/emission line is found, its width (disk centre), features to be observed in that emission and its height above the photosphere. This data is by no means complete and will be updated if the data is found.
Table 1 The most common elements found in the Sun and the wavelength (nm) where the absorption line is found within the spectrum, and the features to be observed in the emission.
Figure 3 shows the most famous chart plotted in 1981 by Vernazzi et al. indicating the height above the photosphere where these lines occur. It is the most cited chart in the literature regarding line heights in the solar chromosphere.
Figure 3 Line heights from 'The Solar Chromosphere' Vernazzi et al. 1981 The Astrophysical Journal Supplement series.
The most common emission lines used by amateur astronomers to view the chromosphere are Calcium II K and Hydrogen alpha. Below (figure 4) are the absorption lines in the spectrum of Calcium II K and H showing the different regions where we should be observing the emission to see maximum line height details.
Figure 4 The regions of Calcium II K and H where different heights above the photosphere can be observed.
As you can see in Calcium II K the central region is called K3 and the two smaller absorption peaks are the K1 regions. This is exactly mirrored in the Calcium II H. The two Calcium lines should be thought of as a doublet like the sodium doublet, with the only difference being how the transitions are made with the two outer shell electrons.
Most amateurs will discuss the bandwidth in which a filter performs, with the 'narrower the better'. Figure 5 demonstrates the bandwidth of some of the common commercial filters.
Figure 5 Some common commercial filter bandwidths overlaid on the Calcium II K and H and Hydrogen alpha absorption spectra.
What you will see is the narrower the bandwidth of the filter, the more likely you will be able to observe the faint emission line height details without leakage from the photosphere (which is much brighter). This is more complex than it seems and will be discussed in part 2 http://solarnutcase.livejournal.com/9896.html
For example in Calcium II K a Coronado PST is 2.2A, a Lunt Solar system filter is 2.4A, whereas a Baader K line filter is up to 80A (too large for this chart). For an Omega Calcium H filter or a Daystar filter system at 5A, this is too large to see subtle line height detail as too much of the photospheric leakage will be included. The Calcium II H line is also very close to the hydrogen epsilon line which could interfere if the filter is not narrow enough.
Less than 2A would be most beneficial in these two Calcium lines. Most professional telescope systems use increments of 0.2A to be able to observe in the separate K1, K2 and K3 regions of Calcium II K. This can only be replicated by the amateur using a spectrohelioscope.
1 angstrom = 0.1nm