Stellar Spectra and Classification
A star's spectrum is like a fingerprint: the pattern of absorption lines in its light tells you its temperature, composition, and spectral type. By analyzing which wavelengths get absorbed, astronomers can classify stars from the hottest O-types to the coolest Y-type brown dwarfs.
Radiation and Stellar Spectra
Stars emit blackbody radiation, which is a smooth, continuous spectrum of light across all wavelengths. The peak wavelength of that emission depends on temperature. Wien's displacement law describes this relationship: hotter stars peak at shorter (bluer) wavelengths, and cooler stars peak at longer (redder) wavelengths.
But we don't just see a smooth continuum when we observe a star. The star's atmosphere (the outer layer of gas) modifies the spectrum:
- Atoms and molecules in the cooler outer layers absorb light at specific wavelengths, creating dark absorption lines in the continuous spectrum.
- In certain conditions (very hot or active stars, for example), emission lines can appear instead, where gas emits light at specific wavelengths.
- The opacity of the atmosphere determines which wavelengths of light can escape, and the degree of ionization in hot atmospheres shapes which lines are present.
The result is a continuous spectrum with a characteristic set of absorption (or sometimes emission) lines superimposed on it. That pattern of lines is what astronomers use to classify stars.
Temperature Effects on Absorption Lines
A star's surface temperature is the main factor controlling which absorption lines show up in its spectrum. Here's why:
Temperature determines the energy states of atoms in the star's outer layers. At different temperatures, atoms occupy different energy levels, and only atoms in the right energy state can absorb a given wavelength of light.
- Cooler stars (like red dwarfs) have more absorption lines in the visible spectrum. Their lower temperatures keep atoms in lower energy states, so many different transitions are possible. Molecules can even survive intact, adding molecular absorption bands.
- Hotter stars (like blue giants) have fewer visible absorption lines. High temperatures excite or ionize atoms to higher energy states, which reduces the number of atoms available to absorb at visible wavelengths.
The strongest absorption lines correspond to the most abundant elements at the right energy state for a given temperature:
- Hydrogen lines are most prominent around 10,000 K (spectral type A), where hydrogen atoms are in the right energy state for Balmer-series absorption.
- Calcium lines (Ca II) are strongest around 6,000 K (spectral type G), not because G-type stars have more calcium, but because the temperature puts calcium atoms in the ideal state for absorption.
This is a key point: line strength depends on temperature, not just abundance. A star might be loaded with hydrogen, but if it's too hot or too cool, the hydrogen lines will be weak.
Characteristics of Spectral Classes
Stars are classified into spectral types based on their surface temperature and the features in their spectra. The classes, arranged from hottest to coolest, are: O, B, A, F, G, K, M, L, T, Y.
A classic mnemonic is "Oh Be A Fine Guy/Girl, Kiss Me Lovingly, Thank You."
Each class has a characteristic temperature range, color, and set of spectral features:
| Class | Temperature (K) | Color | Key Spectral Features |
|---|---|---|---|
| O | > 30,000 | Blue | Ionized helium (He II), highly ionized metals; very few lines overall |
| B | 10,000–30,000 | Blue-white | Strong hydrogen Balmer lines, neutral helium (He I) |
| A | 7,500–10,000 | White | Strongest hydrogen Balmer lines of any class |
| F | 6,000–7,500 | Yellow-white | Weaker hydrogen lines, stronger calcium (Ca II) lines |
| G | 5,000–6,000 | Yellow | Strong Ca II lines, many metal lines (the Sun is a G-type star) |
| K | 3,500–5,000 | Orange | Strong metal lines, weak hydrogen lines |
| M | < 3,500 | Red | Molecular bands (especially titanium oxide, TiO); coolest main-sequence stars |
| L | 1,300–2,500 | Dark red | Metal hydride bands, alkali metal lines (cool low-mass stars and brown dwarfs) |
| T | 700–1,300 | — | Methane absorption bands dominate (cool brown dwarfs) |
| Y | < 700 | — | Ammonia absorption features (the coolest known brown dwarfs) |
Notice the trend: as you move from O to M, spectra go from showing ionized atoms with few lines to showing neutral atoms and even molecules with many lines. The L, T, and Y classes extend this into the realm of brown dwarfs, where temperatures are low enough for complex molecules like methane and ammonia to form.
Brown Dwarfs vs. Planets
Brown dwarfs occupy a gray zone between stars and planets. The distinction comes down to mass and the ability to sustain nuclear fusion.
- Brown dwarfs have masses between roughly 13 and 80 Jupiter masses. They're massive enough to fuse deuterium (a heavy form of hydrogen) in their cores, but not massive enough to sustain the regular hydrogen fusion that powers true stars. Deuterium fusion provides energy temporarily, but once the deuterium runs out, the brown dwarf simply cools and contracts over time.
- Planets have masses below about 13 Jupiter masses. They lack the mass to sustain any type of core fusion and form through accretion of material in a protoplanetary disk around a young star.
The boundary between the two is genuinely blurry. Objects near 13 Jupiter masses are sometimes called "sub-brown dwarfs" or "super-Jupiters," and there's ongoing debate about whether mass alone or formation history should define the category.
To make things trickier, brown dwarfs and planets can have overlapping temperatures. The coolest Y-type brown dwarfs have surface temperatures comparable to some planetary atmospheres (a few hundred kelvin). At that point, you can't tell them apart by temperature or spectrum alone. Distinguishing them requires either measuring the object's mass or figuring out how it formed.