Since the discovery of the chromosphere and since the hand-drawings of Secchi (1877) we have been able to observe this solar atmospheric layer in much detail. Many theoretical models have been proposed to understand its peculiar characteristics. But, only in the last recent years we have been able to address the problem with fine spectropolarimetry and high spatial resolution. We can study the fine details and resolve small structures, following their dynamics in time. Within these recent advances it has been possible both to test current theories and to observe new unexpected phenomena. This work thus aims at contributing to the understanding of the solar chromosphere.
This first Chapter provided a broad introduction to the context of this work. We have briefly presented some general properties of the Sun and the chromosphere. In the fol-lowing pages, throughout Chapter 2, we summarize some theoretical concepts of radiative transfer and spectral line formation needed for this work. We also present general char-acteristics of the two spectral lines studied: Hαand He10830 Å. Chapter 3 presents in detail the observations. There we also summarize the characteristics of the used telescope and optical instruments, as well as the data reduction and post-processing methods ap-plied to achieve spatial resolutions better than 0.005. Next, in Chapter 4, we discuss results from data on the solar disc, dealing with the chromospheric dynamics and fast events ob-served in our data. We present the observations of magnetoacustic waves as well as other fast events. Chapter 5 is devoted to the spicules above the solar limb. The analysis of the spectroscopic intensity profiles from spicules in the infrared spectral range can be used to compare current theoretical models with observations. Further, we present high resolution images in Hαof spicules. Finally, the concluding Chapter 6 of this thesis summarizes the main conclusions and gives an outlook for future work.
2It is very common in astrophysics, specially in solar physics, to use magnetic field strength synony-mously with magnetic flux density. The reason is that in most astrophysical plasmas B=H in Gaussian units.
We follow this use in this thesis.
Most of the information from the extraterrestrial cosmos, also from the Sun, arrives as radiation from the sky. It comes encoded in the dependence of the intensity on direction, time and wavelength. Also, the polarization state of the light contains information. These characteristics of the light we observe from any object have their origin in the interaction of atoms and photons under the local properties (temperature, density, magnetic field, radiation field itself, . . . ).
To extract this encoded information from the recorded intensities it is important to un-derstand how the radiation is created and transported in the cosmic plasmas and released into the almost empty space.
This Chapter describes in the following sections the basis of radiative transfer and spectral line formation. We continue discussing the special properties of the spectral lines used in this work: the hydrogen Balmer-αline (named Hαfor short) at 6563 Å, and the He10830 Å multiplet.
2.1 Radiative transfer and spectral line formation
Light, consisting of photons, interacts with the gas (of the solar atmosphere, in our case) via absorption and emission. Let Iλ(~r,t, ~Ω) be the specific intensity (irradiance) at the point~rin the atmosphere, at timet, and into directionΩ~, with|Ω~|=1. We further denote byκλandλas the absorption and emission coefficients, respectively.
Along a distance dsin the directionΩ~, the change ofIλ is given by
dIλ =−κλIλds+λds, (2.1)
or dIλ
ds =−κλIλ+λ. (2.2)
We define also the optical thickness between some points 1 and 2 in the atmosphere by
dτλ = −κλds ; τλ,1−τλ,2 =− Z 1
2
κλds, (2.3)
and the source functionSλof the radiation field as Sλ = λ
κλ . (2.4)
In the solar atmosphere, absorption and emission are usually effected by transitions between atomic or molecular energy levels, i.e. by bound-bound, bound-free and free-free transitions. If collisions among atoms and with electrons occur much more often than the radiative processes, the atmospheric gas attains statistical thermal properties such as Maxwellian velocity distributions and the population and ionization ratios according to the Boltzamnn and Saha formulae. These properties define locally a temperatureT. It can be shown (e.g. Chandrasekhar 1960) that in these cases, calledLocal Thermodynamic Equilibrium (LTE), the source function is given by the Planck function or black body radiation
Sλ = Bλ = 2hc2 λ5
1
ehc/λkT −1. (2.5)
Sλ varies much more slowly with wavelength than the absorption/emission coefficients across a spectral line. Thus, within a spectral line,Sλ can be considered independent of λ.
Generally, LTE does not hold, especially in regions with low densities (thus with only few collisions relative to radiation processes) and near the outer boundary of the atmosphere from where the radiation can escape into space. The solar chromosphere is a typical atmospheric layer where non-LTE applies. In this case, the population densities of the atomic levels for a specific transition depend on the detailed processes and routes leading to the involved levels.
Equation 2.2 has the following formal solution I(τ2)= I(τ1)e−(τ1−τ2)+Z τ1 A second order expansion ofS(τλ) leads to the Eddington-Barbier relation
Iλ(τλ =0)≈ Sλ(τλ =1). (2.8) This says that the observed intensity Iλ at a wavelength λis approximately given by the source function at optical depthτλ =1 at this same wavelength. In LTE, the intensity then follows the Planck functionBλ(T(τλ = 1)).
In spectral lines, the opacity is much increased over the continuum opacity. Since the temperature decreases with height in the solar photosphere the intensity in spectral lines is decreased, and we probe higher and cooler layers. This explains the formation of absorption lines in LTE.
In non-LTE, when collisional transitions between atomic levels occur seldom and near the outer atmospheric border, photons can escape and are thus lost for the build-up of a radiation field in the specific transition. Then, the upper level of the transition becomes underpopulated and the source function has decreased below the Planck function at the local temperature. It follows that, even for constant temperature atmospheres, a strong
absorption line can be observed.
Outside the solar limb, in the visible spectral range, one observes spectral lines (and very weak continua) in emission. In spectral lines, high chromospheric structures are seen in front of a dark background.