5.5 Formation of Spectral Lines

Formation of Spectral Lines

Spectral lines are the fingerprints of atoms. When light from a star or nebula is spread into a spectrum, bright or dark lines appear at very specific wavelengths. These lines tell astronomers what elements are present, how hot the gas is, and even how ionized it is. This section covers how emission and absorption lines form, what ions are, and how spectral lines act as a thermometer for distant objects.

Formation of Spectral Lines

Formation processes of emission and absorption line spectra

Emission and absorption lines are two sides of the same coin. Both arise from electrons jumping between specific energy levels inside atoms, but the circumstances that produce each type are different.

Emission line spectra form in a hot, low-density gas. Here's the process:

1. An energy source (electric current, nearby hot star, shock wave) excites atoms in the gas.
2. Electrons in those atoms jump up to higher energy levels.
3. The electrons are unstable at those higher levels, so they quickly fall back down.
4. As each electron drops, it emits a photon whose wavelength corresponds exactly to the energy difference between the two levels.
5. The result is a set of bright lines at specific wavelengths against a dark background.

A common example: hydrogen gas excited by an electric current in a discharge tube produces the bright red, blue-green, and violet lines of the Balmer series.

Absorption line spectra form when a continuous spectrum passes through a cooler, low-density gas:

1. A hot, dense source (like the interior of a star) emits a continuous spectrum containing all wavelengths.
2. That light passes through a cooler gas (like the star's outer atmosphere).
3. Atoms in the cooler gas absorb photons at the exact wavelengths that match their electron energy transitions.
4. Those absorbed wavelengths are removed from the continuous spectrum, leaving dark lines.
5. The dark lines appear at the same wavelengths where emission lines would appear if that same gas were excited.

The most famous example is the Fraunhofer lines in the solar spectrum. Hundreds of dark lines appear because the Sun's cooler outer layers absorb specific wavelengths from the continuous light produced deeper inside.

The key link: emission and absorption lines for a given element occur at identical wavelengths. Whether you see bright lines or dark lines depends on the geometry, specifically whether you're looking at the hot gas itself or through a cooler gas in front of a brighter source.

Ions in astronomical contexts

An ion is an atom that has lost or gained one or more electrons, giving it a net electric charge. In astronomy, you'll almost always deal with atoms that have lost electrons (positive ions, or cations), because the environments are so energetic.

Ionization happens when an atom absorbs enough energy to knock one or more electrons free. Two main mechanisms drive this in space:

• Photoionization: High-energy photons (ultraviolet, X-ray, or gamma-ray) strike an atom and eject an electron. This is common near hot stars and in supernova remnants.
• Collisional ionization: In extremely hot environments, atoms move fast enough that collisions between particles can knock electrons loose. This occurs in stellar cores, accretion disks around black holes, and the solar corona.

The degree of ionization in a gas depends on two factors:

• Temperature: Higher temperatures mean atoms have more kinetic energy, making it easier to strip away electrons. Hot stellar cores are highly ionized.
• Density: Lower densities actually favor ionization, because there are fewer opportunities for free electrons to recombine with ions. The thin gas of the interstellar medium can stay ionized more easily than a dense gas at the same temperature.

Astronomers use a notation system for ionization states. For example, Fe I means neutral iron, Fe II means singly ionized iron (one electron removed), Fe III means doubly ionized, and so on.

Spectral indicators of gas temperature

Different ionization states produce different spectral lines, so the pattern of lines you observe acts as a thermometer for the gas.

Cooler gases tend to have atoms in lower ionization states (neutral or singly ionized). Spectral lines from these species are stronger in cooler objects like red giants and molecular clouds.

Hotter gases strip away more electrons, producing higher ionization states (doubly, triply ionized, or more). Lines from these species dominate in the spectra of blue supergiants and active galactic nuclei.

You can estimate gas temperature by comparing the relative strengths of lines from different ionization levels:

1. If lines from both neutral and singly ionized atoms are present with similar strength, the gas is at an intermediate temperature. The solar photosphere (about 5,800 K) is a good example.
2. If lines from highly ionized atoms dominate and neutral lines are weak or absent, the gas is very hot. The solar corona (over 1,000,000 K) shows this pattern.

Certain lines also serve as temperature thresholds. For instance, He II lines (from singly ionized helium) require temperatures above roughly 50,000 K to form. If you see He II lines in a spectrum, you immediately know the gas is at least that hot. This makes specific spectral lines useful as minimum-temperature indicators.

Fundamental concepts in spectral line formation

Several underlying ideas tie this topic together:

• Quantum mechanics explains why energy levels in atoms are discrete (not continuous). Electrons can only occupy specific energy levels, which is why spectral lines appear at precise wavelengths rather than smeared across the spectrum.
• Blackbody radiation provides the continuous spectrum background. A hot, dense object emits light at all wavelengths, and this continuous spectrum is what gets "stamped" with absorption lines when it passes through cooler gas.
• Atomic excitation is the process of an electron moving to a higher energy level after absorbing energy. This is the starting point for both emission (the electron later falls back down and emits a photon) and absorption (the photon that caused the jump is removed from the spectrum).
• Spectroscopy, the technique of splitting light into its component wavelengths and analyzing the resulting spectrum, is the primary tool astronomers use to determine the composition, temperature, density, and motion of celestial objects.