Chapter 1 General Introduction

1. Different ionized gaes

• H II regions: to probe the evolution of the elements and the star formation history of galaxies
  • There are O- or early B-type stars (young, hot, luminous stars) with effective temperature of $3\times {10}^4\mathrm{K}-5\times {10}^4\mathrm{K}$.
  • H is ionized, part of the He is ionized.
  • Typical density: ${10}^{1-4}{\mathrm{cm}}^{-3}$
  • Velocity of gas $\approx 10\mathrm{km}/\mathrm{s}$, close to the isothermal sound speed.
  • Strong H I recombination lines, [N II] and [O II] collionally excited lines.
  • Concentrated to the spiral arms for the Milky Way.
• Planetary nebulae: the outer remaining envelopes of dying stars
  • There are dying stars with effective temperature of $\sim 5\times {10}^4\mathrm{K}$.
  • More ionized than H II region
  • Typical density: ${10}^{2-4}{\mathrm{cm}}^{-3}$
• Supernova remnants: to observe material from the burned-out deep interiors of supernovae
• Gas around active galactic nuclei: to study the central engines of galaxies

2. Spectra

Figure 1.1 in this book

• Emission lines: dominated by forbidden lines like [O III] $\lambda \lambda 4959,5007$, [NII] $\lambda \lambda 6548,6583$, and [S II] $\lambda \lambda 9069,9523$; by permitted lines like H$\alpha$ and H$\beta$. More details are available in Figure 1.1.
• Weak continuous spectra: consist of atomic and reflection components.
  • Atomic continuum: chiefly by free-bound transitions, mainly in the Paschen continuum of H I at $\lambda >3646\mathring{\mathrm{A}}$, and the Balmer continuum at $912\mathring{\mathrm{A}}<\lambda <3646\mathring{\mathrm{A}}$.
  • Reflection continua: starlight scattered by dust.
• Continous spectrum is strong in the radio-frequency region.

3. Physical conditions

• The energy source: ultraviolet radiation from stars. The effective surface temperature of hot stars can be as high as $\gtrsim 3\times {10}^4\mathrm{K}$.
• The main energy-input mechaism: the photoionization of H.
• The liberated photoelectrons collide with ions and hence excite the low energy levels of the ions.
• Although the probabilities of the radiative decay of the excited levels is small, the collisional de-excitation is even more inefficient as a consequence of low density ($n_e\le {10}^4{\mathrm{cm}}^{-3}$). Therefore, almost every excitation leads to emission of a photon.
• There is an equilibrium between the photoionization and recombination processes, which is called the ionization equilibrium.
• The recombination of a ion will form a excited atoms, such as H${}^{+}$ give rise to excited atoms of H and leads to the emission of H I (the origin of the H I Balmer- and Paschen-line spectra).
• Inelastic collisions of free electrons and ions converts kinetic energy into excitation energy, which is called the collisional excitation. The inverse process is called the collisional de-excitation. There are two ways to lose the excitation energy for a excited ions: (1) radiative decay; (2) collisional de-excitation.
• There are almost no lines for the ions that are fully ionized in the hot gas. Although a single nucleus can be excited, but the energy is too high, which is $\gtrsim 1000\mathrm{keV}$. NuDat 3.0 gives the energy levels for the nuclei. NuDat 3.0 for ${}^{56}$Fe shows the energy levels for ${}^{56}$Fe. The energy of the first excited level for ${}^{56}$Fe is $0.847\mathrm{MeV}$, which is too high to be excited in the hot gas. The hydrogen nucleus is even more stable.