Saturday, July 15, 2006

Astronomical Data – 3c. Spectroscopy

In today’s post, we’ll be exploring another astronomical technique known as spectroscopy in which astronomers can use light to determine chemical composition.

If you’ve forgotten where light comes from, you might want to go back and reread my earlier post on where light comes from since we’ll be referring back to that quite a bit in this post. Assuming you do remember, let’s take an expanded look at all that.

You should recall that we can get the entire continuous spectrum of light (ie, the entire rainbow including what lies beyond just the visible part) due to electrons falling into orbitals from outside the atom. Since they can fall any distance, that’s what makes it possible to have every wavelength thus making the spectrum complete.

However, in this scenario I made a few assumptions that you may or may not have caught.

The first is that there are slots for the electrons to fall into. If the electron is already neutral (ie, has as many electrons as there are protons in the nucleus) then the atom won’t readily accept any more.

The second is that, if an atom is needing an electron to become neutral, that there’s actually some nearby to fall in.

On the surface of stars, this isn’t a bad assumption. The heat causes atoms to become ionized (lose their electrons) and since the star is relatively dense, those electrons are right nearby to fall into another atom, give off a photon, and then get ionized again.

But in case you haven’t noticed, stars aren’t the only thing in the galaxy. So, as you might be anticipating, continuous spectrums aren’t the only kind.

Let’s consider another case in which we have a low density gas that’s warm but not warm enough to actually ionize the atom (for the time being, stick with the hydrogen model in your mind since it’s easiest). In this, since the atom will be imparted with some energy, but not enough to ionize it, the electrons will jump up orbitals and fall back down, emitting a photon.

But since, as I discussed in the previous post, the orbitals have discreet energies, the photons emitted will be confined to those energies as well, and thus, only certain wavelengths. So for hydrogen, here’s what that looks like in the visible part of the spectrum:

Since hydrogen is the most plentiful element in the universe, we see this pattern popping up almost everywhere. However, what happens if we have other gasses in the same condition?

In that case, we get different lines being emitted due to different atomic configuration. Here’s a few more:

As you can see each element has a unique configuration of lines. I generally compare this to a barcode that’s unique to each different element. Thus, if astronomers can put the light from one of these low density, warm clouds through a spectrum and see a pattern that conforms to one of these, they’ll know what element is present. This isn’t always as easy as it sounds for many reasons. One of the major reasons is that there’s frequently more than one element present in what we’re wanting to look at, so sorting things out can get difficult. Another (which we’ll explore later) is that such lines aren’t always where they’re supposed to be due to a few different reasons.

But assuming that things can be sorted out, an accurate determination of chemical makeup can be determined! This type of spectra, one with distinct lines, is called an emission spectra.

So keeping all that in mind, let’s take a look at another scenario.

In this one, let’s allow the cloud of gas to be nice and cold heat isn’t causing any electrons to jump up into higher orbitals. In this case, density is unimportant since the atoms already have their lower energy levels filled and aren’t accepting electrons.

But now let’s imagine that we put a source that emits a continuous spectrum behind it so this cloud is in the way. In this scenario, light from the continuous spectrum source will have to pass through the cold cloud. Here’s where something special happens.

If a photon happens hit an atom of the cold gas, it can get absorbed. But only if it’s of one of the specific energies that corresponds to one of those jumps between orbitals. If that happens, and the photon is absorbed, then the electron will take that energy and hop up.

Of course, since that’s a higher energy level, it will fall back down, emitting a photon of the exact same energy as the one it absorbed.

So what? Photon absorbed and given right back off. What’s the big deal?

The trick here is that, the photon that is given off can be given off in any direction. The chance that it will happen to go straight towards the observer is pretty slim. Thus, the observer no longer sees light at the wavelength that corresponds to that energy!

So now instead of having a continuous spectrum, the observer will see one that has lines subtracted from it at the same wavelengths that the gas the light was passing through would emit if it were hot.

Where do we see this? The major place is in stars.

Yeah yeah, I know I said stars have continuous spectrums earlier. And they would. If it weren’t for the fact that they don’t have solid surfaces and just slowly fade into a sort of extended atmosphere. That atmosphere is (relatively) cool, and thus, will absorb pieces of the continuous spectrum generated lower in the star.

So what’s this called? As you might expect, it’s called an absorption spectra.

Let’s take a look at one.

Want to take a guess what star this is?

It’s the Sun! And wow is that a lot of absorption lines! Some of the most prominent ones are caused by hydrogen and helium. Some of the other, fainter ones are caused by trace gasses in the Sun’s atmosphere, but most are caused by our own atmosphere.

As a brief aside, you’ll also notice that the spectrum is brightest in that yellow green area as I pointed out in a few other posts.

Incidentally, the late mid 1800’s was the first time the solar spectrum was examined. At that time, astronomers first determined that the sun was made mostly of hydrogen and were able to pick out many other elements. However, some prominent lines couldn’t be explained. Thus, the presence of a yet undiscovered element was inferred. This element was named Helium after the Greek word for sun, Helios. Helium was later discovered on Earth.

So let’s do some recapping before we go any further:

We’ve now looked at three types of spectra: continuous, emission, and absorption.

With a continuous spectrum or absorption spectra, we’ll be able to find the wavelength where the most light is given off (as I pointed out in my last post), which can give us the temperature of the star.

Emission spectra and absorption spectra are useful because the pattern of lines tell us what chemicals are present.

Continuous spectra come from hot, high density gasses.

Emission spectra come from low density, warm gas.

Absorption spectra come from continuous spectra passing through cool gas.

But the fun doesn’t end there!

Let’s take another look at that absorption spectra. But this time, let’s do it in a graphical form:

The blue line here is what the continuous spectrum would look like for this star if there weren’t the deep absorption lines present. Those very deep ones at ~430, ~480, ~520, and ~655 nm are those caused by hydrogen we looked at earlier.

You’ll notice they’re pretty damned deep. However, the other ones aren’t so deep. The reason has to do with the abundance of each element in that star’s atmosphere. Since there’s a lot of hydrogen in stars, it makes sense that the hydrogen lines be the deepest.

So by looking at how deep each line is astronomers can figure out the ratio of elements which is a pretty nifty trick.

These absorption lines can also be used in other ways. Another use is that they reveal the presence of magnetic fields thanks to an effect known as the Zeeman effect which causes the spectral lines to split into two if there’s a magnetic field present.

If you recall the solar telescope at Mt. Wilson I talked about, it looks for this spectral line splitting at thousands of different points on the face of the sun, which allows the astronomers working on that project to essentially map the magnetic field.

So that’s it for this post. In my next post, we’ll look at some other uses of these lines to determine other quantities, but since it will require another bout of background explanation, I’ll save that for the next post.

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