KAHIBARO
Table of Contents
Light from Stars
Stellar spectra are one of the most powerful tools in astrophysics. A star is far away, and we usually cannot touch it, sample it, or visit it. But starlight carries information. By spreading that light into its different wavelengths, astronomers can learn what a star is made of, how hot it is, how fast it is moving, and even some details about its surface conditions.
A stellar spectrum is a record of how much light a star emits at different wavelengths. Instead of seeing only a single bright point, we separate the light into colors, from shorter wavelengths to longer wavelengths. The result is not always a smooth rainbow. It often contains dark or bright lines, and these features are the key to understanding stars.
Continuous Spectra and Spectral Lines
A hot dense object produces a continuous spectrum, meaning it emits light over a broad range of wavelengths. A star’s inner layers behave approximately like this. But the outer layers of a star are cooler and thinner. As light passes through them, atoms and ions absorb certain wavelengths. This creates dark absorption lines in the spectrum.
Sometimes gases can also produce bright emission lines, but for most ordinary stars the dominant visible feature is a continuous background crossed by absorption lines.
The basic idea is simple. Each chemical element interacts with light in its own way. Hydrogen has one pattern of lines, helium has another, sodium another, and so on. By matching the lines seen in starlight with patterns measured in laboratories on Earth, astronomers identify the elements present in stars.
A stellar spectrum is not just a rainbow. It is a graph or image showing intensity as a function of wavelength, and its spectral lines reveal the physical properties of the star.
Why Spectral Lines Appear
Atoms have allowed energy levels. An electron in an atom can move between these levels only by gaining or losing specific amounts of energy. Light comes in photons, each with energy
$$
E = hf = \frac{hc}{\lambda},
$$
where $h$ is Planck’s constant, $f$ is frequency, $c$ is the speed of light, and $\lambda$ is wavelength.
If a photon has exactly the right energy, an atom can absorb it. That removes light at that wavelength from the star’s spectrum, producing an absorption line. Because every element has its own set of energy levels, every element has its own spectral fingerprint.
This is why spectra are so useful. A star that looks white to the eye may actually contain a rich pattern of many lines.
What Stellar Spectra Tell Us
From stellar spectra, astronomers can determine several important properties.
The first is chemical composition. If the spectrum contains hydrogen Balmer lines, calcium lines, sodium lines, or iron lines, those elements are present in the star’s atmosphere.
The second is temperature. The strength of particular lines depends strongly on temperature. A very hot star may show strong ionized helium lines, while a cooler star may show strong neutral metal lines or molecular bands.
The third is motion along the line of sight. If the lines are shifted from their usual wavelengths, the star is moving toward us or away from us. This is the Doppler effect.
The fourth is surface gravity and pressure. Dense stellar atmospheres can broaden lines. In this way, line shape gives clues about the physical environment in the star’s outer layers.
The fifth is rotation. If one side of a rotating star moves toward us while the other side moves away, spectral lines become broadened. Faster rotation usually means broader lines.
Spectral Classification
Astronomers group stars into spectral classes based on the appearance of their spectra. The main sequence of classes is
$$
O, B, A, F, G, K, M
$$
from hottest to coolest.
This ordering was originally based on observed spectral features, but it also corresponds to temperature. Each class has characteristic lines and colors.
| Spectral Class | Approximate Temperature | Typical Appearance in Spectrum |
|---|---|---|
| O | above $30000 \, \text{K}$ | ionized helium, weak hydrogen |
| B | $10000$ to $30000 \, \text{K}$ | helium, stronger hydrogen |
| A | $7500$ to $10000 \, \text{K}$ | strongest hydrogen lines |
| F | $6000$ to $7500 \, \text{K}$ | hydrogen weaker, metals stronger |
| G | $5200$ to $6000 \, \text{K}$ | metals and ionized calcium, like the Sun |
| K | $3700$ to $5200 \, \text{K}$ | strong metal lines, cooler appearance |
| M | below $3700 \, \text{K}$ | molecular bands, especially titanium oxide |
The Sun is a G-type star, more precisely about G2.
A famous memory aid is, "Oh Be A Fine Girl, Kiss Me," though the words are only a mnemonic and not physics.
The spectral sequence $OBAFGKM$ runs from hottest stars to coolest stars, not from brightest to dimmest.
Temperature and Line Strength
It may seem surprising that hydrogen, the most abundant element in many stars, does not always produce the strongest visible lines. The reason is that line strength depends not only on how much of an element is present, but also on the physical state of its atoms.
At low temperatures, many electrons stay in lower energy states and may not be able to absorb visible photons in the relevant way. At very high temperatures, atoms may become ionized, so the needed electrons are no longer bound in the same form. At intermediate temperatures, the visible hydrogen Balmer lines become strongest. This is why A-type stars show especially strong hydrogen lines.
So a weak hydrogen line does not necessarily mean little hydrogen. It may simply mean the temperature is not right for that line to appear strongly.
Absorption Spectra of Stars
Most ordinary stellar spectra are absorption spectra. A hot interior produces broad continuous light, and cooler gas above it absorbs selected wavelengths.
A simple picture is shown below.
This picture is simplified, but it captures the main idea.
Doppler Shift in Stellar Spectra
If a star moves relative to Earth, its spectral lines shift in wavelength. If the star moves away, the lines shift toward longer wavelengths, called redshift. If the star moves toward us, the lines shift toward shorter wavelengths, called blueshift.
For speeds much smaller than the speed of light, the radial velocity relation is approximately
$$
\frac{\Delta \lambda}{\lambda} \approx \frac{v}{c},
$$
where $\Delta \lambda$ is the change in wavelength, $\lambda$ is the original wavelength, $v$ is the speed along the line of sight, and $c$ is the speed of light.
This method lets astronomers measure stellar motion even when the motion is too small to see directly in the sky.
If spectral lines are shifted, compare them with their known laboratory wavelengths. For small speeds,
$$
\frac{\Delta \lambda}{\lambda} \approx \frac{v}{c}.
$$
A positive shift toward longer wavelength means the star is receding.
Line Broadening
Real spectral lines are not infinitely thin. They have some width. Several effects can broaden them.
Thermal motion in the stellar atmosphere causes atoms to move randomly, which creates Doppler broadening.
Pressure in denser atmospheres causes atoms to interact more strongly, which can broaden lines.
Rotation can broaden lines because one edge of the star moves toward us while the opposite edge moves away from us.
Magnetic fields can split or broaden lines, a phenomenon related to the Zeeman effect.
These details help astronomers go beyond simple identification of elements and study physical conditions in stellar atmospheres.
Reading a Simple Stellar Spectrum
A stellar spectrum can be displayed as a graph of intensity versus wavelength. Dark absorption lines appear as dips.
The smooth curve shows the overall distribution of light with wavelength. The narrow dips show where atoms absorbed specific wavelengths.
Spectroscopy as a Tool in Astronomy
The study of spectra is called spectroscopy. It is central to astronomy because stars are too distant for direct sampling. Spectroscopy turns light into information.
Using stellar spectra, astronomers learned that stars contain the same elements found on Earth. They also discovered that the Sun is rich in hydrogen and helium, and they classified stars into temperature groups long before modern space telescopes existed.
This makes stellar spectra one of the foundations of modern astrophysics.
Key Ideas to Remember
A stellar spectrum is the distribution of a star’s light across wavelength. Most stellar spectra show a continuous background crossed by dark absorption lines. These lines occur because atoms in the star’s atmosphere absorb light at specific wavelengths. Each element has a unique pattern of spectral lines, so spectra reveal composition. The pattern and strength of lines also reveal temperature. Shifts in line position reveal motion through the Doppler effect, and line width can reveal pressure, temperature, rotation, or magnetic effects.
Important facts about stellar spectra:
$$
E = hf = \frac{hc}{\lambda}
$$
Atoms absorb and emit only specific photon energies, producing spectral lines.
The main spectral classes are
$$
O, B, A, F, G, K, M
$$
from hottest to coolest.
For small Doppler shifts,
$$
\frac{\Delta \lambda}{\lambda} \approx \frac{v}{c}.
$$
Spectra let astronomers determine composition, temperature, and radial motion of stars.