Astronomy

Why does hydrogen ionization happen in HII regions?

Why does hydrogen ionization happen in HII regions?


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Why does hydrogen ionization happen in HII regions? Why is the hydrogen there ionized?


Stars are responsible.

HII regions$^dagger$ can refer to several things, but usually I guess one thinks of the volumes around star-forming regions. The more massive a star is, the faster it burns its fuel, and at a higher temperature, meaning that the peak of their spectra are more toward the high frequencies. The most massive stars of a stellar population - the so-called O and B stars - produce enough photons with wavelengths below the hydrogen ionization threshold of $lambda = 912$ Å that they carve out bubbles in their surrounding HI clouds, giving rise to HII regions.

Right: The HII region NGC 604 (from Wikipedia). Left: The spectra of three different stars. Only the B star has a significant portion of its spectrum above the hydrogen ionization threshold. Note the logarithmic scale on the intensity (from here).

Because of the high densities, the HII quickly recombines to HI. If the recombination goes directly to the ground state, a new ionizing photon is emitted, which again is absorbed by a hydrogen atom, but if it goes to one of the higher states, the emitted radiation is no longer capable of ionizing the gas. In this way the ionizing radiation is converted into photons of specific wavelengths, corresponding the energy differences between the excited level of hydrogen, most notably the Lyman $alpha$ emission line with $lambda = 1216$ Å.

Because hydrogen is the most abundant element in the Universe, and because Lyman $alpha$ is the most common transition, the Lyman $alpha$ line is an excellent probe of the most distant galaxies where other wavelengths are not observable. This is especially so because the most distant galaxies are also the earliest and hence still in the process of forming, meaning a high star formation rate which, in turn, means that the shortlived OB stars are still present.

In addition to this distinct regions of HII, ionized hydrogen also exist in a more diffuse component between the stars of a galaxy, in huge bubbles caused by stellar feedback and supernovae, and in the intergalactic medium.

$^dagger$The terms HI and HII refers to neutral and ionized hydrogen, respectively.


Why does hydrogen ionization happen in HII regions? - Astronomy

The masters of stellar population synthesis code (building a galaxy from individual stars) are Bruzual & Charlot. The Kennicutt atlas of optical galaxy spectra of various types can be found here. A popular set of UV galaxy templates are those of Coleman, Wu & Weedman. Here is a short introduction to the stellar populations found within typical galaxy spectra.

Galaxy spectra are typically characterized by a strong continuum component, caused by the combination of a range of blackbody emitters spanning a range in temperature. This produces a fairly flat overall spectrum. A primary feature is the strong break at 4000A (Angstroms), caused by the blanket absorption of high energy radiation from metals in stellar atmospheres and by a deficiency of hot, blue stars. A smoothly varying function can easily be fit to spectra, with a clear drop off in intensity at 4000A. Some stellar type spectra show the 4000A break clearly, while others do not, due to the presence of absence of metal lines and a strong blue/UV component. Based on the correlation between the strength of the feature with stellar type, in which type of galaxy would you expect a strong 4000A break - ellipticals , spirals , or irregulars ?

There are absorption features superimposed on the continuum, due to the absorption of atoms (metals) and molecules in stellar atmospheres, and to cold, interstellar gas clouds which siphon off radiation at key frequencies. This implies the presence of old stellar populations, which are typical found in elliptical galaxies and in the bulges of spiral galaxies. Key features include the Calcium H and K lines (found at 3934A and 3969A), the G-band (4304A), and Magnesium (5175A) and Sodium (5894A) lines.

Elliptical Galaxy Spectra
Elliptical galaxy spectra are characterized by strong absorption lines, due to metals in the stellar atmospheres of the low luminosity stellar population. We see few to no emission lines ([OII]3727A and/or [NII]6583A are occasionally present), as there are essentially no young stars and no gas.
[Steward Observatory, R. Kennicutt]

One will also see emission features, due to gas being heated and then re-radiating energy at specific wavelengths. Young stars form within gas clouds, which they then ionize. The emission from the Orion nebula, for example, is fueled by four bright O stars, which emit most of the ionizing photons (E > 13.6 eV) that energize the surrounding HII region.

The center of the Orion Nebula is a churning, turbulent star factory set within a maelstrom of flowing, luminescent gas. Though this 2.5 light-year view is a small portion of the entire nebula, it includes a star cluster and almost all of the light from the bright glowing clouds of gas that make up the nebula. [NASA/HST]

Emission features thus point to very hot gas and OB type stars, from the disks of spiral galaxies and from irregular galaxies. Key features include the [OII] doublet (3737A), [OIII] (4959A and 5007A), and the Balmer series (6563A, 4861A, 4340A, 4103A, . ).

Irregular Galaxy Spectra
Irregular galaxy spectra are characterized by strong emission lines, due to hot young stars and surrounding HII regions.
[Steward Observatory, R. Kennicutt]

Most galaxy spectra are redshifted (spectral features have shifted to longer wavelengths than the rest wavelength values), though some few are blueshifted . This is interpreted similarly to a Doppler shift, and implies that the galaxies are moving away from us (redshifted) or towards us (blueshifted). What is the spatial location of most blueshifted galaxies?


Making it Rain in the Circumgalactic Medium

It can be easy to think of galaxies as islands in the universe, floating around in isolation. However, a galaxy is actually surrounded by a huge sea of low-density gas that extends out to its virial radius and beyond. This gas is known as the circumgalactic medium (CGM), and more and more research is showing that the CGM has a crucial role to play in galaxy evolution. Observing the CGM has proven difficult due to its extremely low density, though, so simulations have played a large role in understanding the physics of this region. In today’s paper, the authors detail the effects of running a CGM simulation with significantly increased resolution, capable of resolving cool gas that precipitates in the CGM and rains down on the galaxy.

What do we know about the CGM?

Residing just outside of the galaxy, the CGM is home to large-scale flows of gas that drive galaxy evolution. These gas flows provide fuel for star formation, regulate the interactions between dark matter halos and the intergalactic medium , and contain the energy, mass, and metals of large outflows from a galaxy. In fact, the CGM is predicted to hold at least as many baryons and heavy elements as galaxies themselves, and most of the metals in the universe are found in the CGM. These metals (meaning anything heavier than hydrogen or helium in astronomy terms), deposited by galactic outflows, serve as the dominant coolant for the CGM. They are capable of radiating energy away more easily than elements like hydrogen, so an increased abundance of metals can lead to cooler gas. Consequently, this influx of metals helps to create two phases of gas: “cool” (10,000 Kelvin) gas composed of neutral hydrogen and other elements in low-energy ionization states, and “hot” (300,000 – 1,000,000 Kelvin) gas that contains oxygen, nitrogen, and neon in high-energy ionization states.

Unfortunately, computational work has chronically underproduced the observed abundances of these ions across redshifts by orders of magnitude. Recent work has shown that AGN feedback can increase the abundances of oxygen and other ions in the hot gas, but the discrepancy remains for hydrogen and other ions in the cool gas. In today’s paper, the authors discuss the effect of increased simulation resolution on these discrepancies.

Resolving the Resolution Issue

Perhaps one reason that simulations struggle to reproduce observations of the CGM lies in their resolution limits. Similar to how using more pixels in a television or computer screen gives a better image, increasing the resolution in a simulation means using more cells or particles to obtain a better physical picture of what is going on. However, each increase in resolution increases the computational cost of the simulation. This means your simulation that took a few days to run could instead take a few months.

Consequently, most simulations of galaxies apply their highest resolution to regions of high density where most of the matter is. This is great for figuring out what happens in the dense disk of a galaxy, but not ideal for studying the low-density CGM. Today’s paper runs simulations that force high resolution upon the CGM, reaching resolutions that are comparable to those normally obtained in the disk of the galaxy. This technique is appropriately named Enhanced Halo Resolution (EHR). Figure 1 shows the resolutions obtained by both a normal cosmological simulation and an EHR one for a region encompassing a galaxy and its surrounding filaments .

Figure 1: Plots of resolution for a traditional (AMR – adaptive mesh refinement) and EHR simulation. Each of these grids is made up of many cells, and spatial resolution refers to the physical length (in kiloparsecs) of the smallest cell that is present in a region. In the left panel, many galaxies are present and a particularly massive galaxy lies at the center. Its virial radius is shown by the dotted white line. Resolution in the CGM is roughly 16 times worse than in the disk of the galaxy. On the right, the EHR simulation enforces high resolution approximately to the virial radius, ensuring that interactions within the CGM are given much more computational attention. Taken from Figure 2 in the paper.

What does this computational cost buy you?

By better resolving the gas in the CGM, the authors note that a number of physical effects present themselves. Firstly, the balance of cool and hot gas is shifted, leaving more cool gas and less hot gas than in simulations with lower resolution. The clouds of cool gas that form are also greater in number and smaller in size. Finally, the amount of neutral hydrogen and other low-energy ions found in the cool gas increases, while the abundances of oxygen, nitrogen, and neon in high-energy ionization states fall due to the decrease in hot gas. Coupled with the aforementioned work on AGN feedback, this can bring simulations closer to the observed abundances for these ions.

Figure 2: A galaxy and the CGM in an AMR simulation and an EHR one. A significant increase in HI (neutral hydrogen) can be seen in the EHR simulation. Recall that neutral hydrogen tracks the cool gas, which condenses into many clumps on the right that weren’t resolved in a traditional AMR simulation. Many of these clumps fall back into the galaxy because they no longer have enough thermal energy to resist the gravitational pull of the galaxy. Taken from Figure 3 in the paper.

In other words, EHR causes more gas in the CGM to cool, condense into clouds, and potentially fall back into the galaxy. This is completely analogous to water vapor in our own atmosphere, which often cools, forms clouds, and rains back down to Earth. In this way, the CGM can be conceptualized as the atmosphere of a galaxy. Figure 2 shows cool gas condensing into these clouds, some of which fall into the galaxy.

Why does an increase in resolution result in more cool gas? The answer lies in how gas mixes in simulations. With lower resolution, clouds of cool gas are typically resolved only by a few cells, inducing artificial mixing between the hot and cool gas. The authors perform a test simulation demonstrating this, shown in Figure 3.

Figure 3: In this test problem, a 4-kiloparsec-wide cloud of cool gas sits in a flow of hot gas for 260 million years. In the low-resolution test, the boundary of the cloud is only resolved by a few cells. This artificially thick boundary means that much of the cool gas quickly mixes with the hot gas and eliminates the HI (neutral hydrogen). In the high-resolution case, the boundary becomes much thinner, allowing the interior cool gas to survive much longer. Taken from Figure 13 in the paper.

Resolution clearly makes a big difference in understanding the physics of the CGM and galaxies. For example, just like plants on Earth sprout after a rain, cool gas that condenses in the CGM and falls into a galaxy can trigger star formation. Understanding the ecology and geology of Earth requires a detailed picture of the atmosphere, and perhaps unlocking the mysteries of galaxy evolution may depend just as strongly on our understanding of the CGM.


Making It Rain in the Circumgalactic Medium

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites the original can be viewed at astrobites.org.

Title: The Impact of Enhanced Halo Resolution on the Simulated Circumgalactic Medium
Authors: Cameron B. Hummels, Britton D. Smith, Philip F. Hopkins, et al.
First Author’s Institution: TAPIR, California Institute of Technology
Status: Submitted to ApJ

It can be easy to think of galaxies as islands in the universe, floating around in isolation. However, a galaxy is actually surrounded by a huge sea of low-density gas that extends out to its virial radius and beyond. This gas is known as the circumgalactic medium (CGM), and more and more research is showing that the CGM has a crucial role to play in galaxy evolution. Observing the CGM has proven difficult due to its extremely low density, though, so simulations have played a large role in understanding the physics of this region. In today’s paper, the authors detail the effects of running a CGM simulation with significantly increased resolution, capable of resolving cool gas that precipitates in the CGM and rains down on the galaxy.

What Do We Know About the CGM?

Residing just outside of the galaxy, the CGM is home to large-scale flows of gas that drive galaxy evolution. These gas flows provide fuel for star formation, regulate the interactions between dark matter halos and the intergalactic medium , and contain the energy, mass, and metals of large outflows from a galaxy. In fact, the CGM is predicted to hold at least as many baryons and heavy elements as galaxies themselves, and most of the metals in the universe are found in the CGM. These metals (meaning anything heavier than hydrogen or helium in astronomy terms), deposited by galactic outflows, serve as the dominant coolant for the CGM. They are capable of radiating energy away more easily than elements like hydrogen, so an increased abundance of metals can lead to cooler gas. Consequently, this influx of metals helps to create two phases of gas: “cool” (10,000 Kelvin) gas composed of neutral hydrogen and other elements in low-energy ionization states, and “hot” (300,000–1,000,000 Kelvin) gas that contains oxygen, nitrogen, and neon in high-energy ionization states.

Unfortunately, computational work has chronically underproduced the observed abundances of these ions across redshifts by orders of magnitude. Recent work has shown that AGN feedback can increase the abundances of oxygen and other ions in the hot gas, but the discrepancy remains for hydrogen and other ions in the cool gas. In today’s paper, the authors discuss the effect of increased simulation resolution on these discrepancies.

Resolving the Resolution Issue

Perhaps one reason that simulations struggle to reproduce observations of the CGM lies in their resolution limits. Similar to how using more pixels in a television or computer screen gives a better image, increasing the resolution in a simulation means using more cells or particles to obtain a better physical picture of what is going on. However, each increase in resolution increases the computational cost of the simulation. This means your simulation that took a few days to run could instead take a few months.

Consequently, most simulations of galaxies apply their highest resolution to regions of high density where most of the matter is. This is great for figuring out what happens in the dense disk of a galaxy, but not ideal for studying the low-density CGM. Today’s paper runs simulations that force high resolution upon the CGM, reaching resolutions that are comparable to those normally obtained in the disk of the galaxy. This technique is appropriately named Enhanced Halo Resolution (EHR). Figure 1 shows the resolutions obtained by both a normal cosmological simulation and an EHR one for a region encompassing a galaxy and its surrounding filaments .

Figure 1: Plots of resolution for a traditional (AMR — adaptive mesh refinement) and EHR simulation. Each of these grids is made up of many cells, and spatial resolution refers to the physical length (in kiloparsecs) of the smallest cell that is present in a region. In the left panel, many galaxies are present and a particularly massive galaxy lies at the center. Its virial radius is shown by the dotted white line. Resolution in the CGM is roughly 16 times worse than in the disk of the galaxy. On the right, the EHR simulation enforces high resolution approximately to the virial radius, ensuring that interactions within the CGM are given much more computational attention. [Hummels et al. 2019]

What Does this Computational Cost Buy You?

By better resolving the gas in the CGM, the authors note that a number of physical effects present themselves. Firstly, the balance of cool and hot gas is shifted, leaving more cool gas and less hot gas than in simulations with lower resolution. The clouds of cool gas that form are also greater in number and smaller in size. Finally, the amount of neutral hydrogen and other low-energy ions found in the cool gas increases, while the abundances of oxygen, nitrogen, and neon in high-energy ionization states fall due to the decrease in hot gas. Coupled with the aforementioned work on AGN feedback, this can bring simulations closer to the observed abundances for these ions.

Figure 2: A galaxy and the CGM in an AMR simulation and an EHR one. A significant increase in HI (neutral hydrogen) can be seen in the EHR simulation. Recall that neutral hydrogen tracks the cool gas, which condenses into many clumps on the right that weren’t resolved in a traditional AMR simulation. Many of these clumps fall back into the galaxy because they no longer have enough thermal energy to resist the gravitational pull of the galaxy. [Hummels et al. 2019]

Why does an increase in resolution result in more cool gas? The answer lies in how gas mixes in simulations. With lower resolution, clouds of cool gas are typically resolved only by a few cells, inducing artificial mixing between the hot and cool gas. The authors perform a test simulation demonstrating this, shown in Figure 3.

Figure 3: In this test problem, a 4-kiloparsec-wide cloud of cool gas sits in a flow of hot gas for 260 million years. In the low-resolution test, the boundary of the cloud is only resolved by a few cells. This artificially thick boundary means that much of the cool gas quickly mixes with the hot gas and eliminates the HI (neutral hydrogen). In the high-resolution case, the boundary becomes much thinner, allowing the interior cool gas to survive much longer. [Hummels et al. 2019]

About the author, Michael Foley:

I’m a graduate student studying Astrophysics at Harvard University. My research focuses on using simulations and observations to study stellar feedback — the effects of the light and matter ejected by stars into their surroundings. I’m interested in learning how these effects can influence further star and galaxy formation and evolution. Outside of research, I’m really passionate about education, music, and free food.


One way of looking at the matter is to realize that the Main Sequence represents a set of stable, equilibrium states for stars as a function of mass. That is, any ball of (mostly) hydrogen that collapses due to its self gravity will continue to collapse until it reaches a state where some means of generating energy becomes available to it. And this energy source supplies the thermal pressure that halts the gravitational collapse.

The existence of a Main Sequence means that the mass of a collapsing protostar will determine the resulting Luminosity and Surface Temperature at which it reaches a stable state.

In the 1860s, William Thompson, Lord Kelvin, and Ludwig von Helmholtz argued that gravitational compression could be the source of the Sun's energy. BUT! It turns out that the predicted lifetime of the Sun would only be about a million years. And it would mean that the Sun is shrinking, and would have been much larger in the past. Even in the 19th century, geologic dating, and the upper limits on measured changes in the Solar radius ruled out the hypothesis that gravitational compression could be the source of Solar power.

Combustion (burning coal, or oil etc.) releases

compared to the Solar luminosity of

This implies the combustion rate would have to be

Now, the Sun contains about atoms, so this would imply that the Solar lifetime is only

Energy generation in Main-Sequence stars is from nuclear fusion. This follows from Albert Einstein's Special Theory of Relativity, and was understood by the 1930s. One of the fundamental results of the Special Theory of Relativity is the equivalence of mass and energy, summed up in the famous equation

Note that the equation implies that a small amount of mass can be converted into a very large amount of energy.

In simplified form, the Sun generates energy by converting four Hydrogen nuclei into a Helium nucleus. If one goes to a table of atomic masses, and adds up the masses involved, one finds that four H atoms are more massive than one He atom by about 0.7%. The mass difference is for a single reaction. This means the energy generated is

Or seven orders of magnitude larger than that generated by combustion. That implies

Actually, only 10% of the Sun's mass will undergo fusion, so the true lifetime is only 10 Billion years, not 100 Billion.


Therefore we can see that stars are just atoms (hydrogen) hitting each other producing light. Scientists know what stars are made of through the use of the a.

(Nuclear Chemistry) The two Hydrogen atoms “combine” to become Helium, a new element. Nuclear Fission occurs when an large unstable element splits into small.

Alpha (α) Radiation The alpha radiations particles have 2 protons and 2 neutrons which are in a helium nucleus. As the alpha particles emit the helium nucle.

c. Neutrinos tells us about the thermonuclear process that happens in the Suns core and what is released from the core, which we cannot learn from other venu.

The chromosphere is characterized by plages and flares. Plages are found in the lower level of the chromosphere nearest the photosphere. Plages are denser wi.

1. What are Kepler’s 3 Laws. Explain each law in detail and why it is important in astronomy. (3 pts) Kepler’s Three Laws that everything orbits the sun. Th.

The Nebular and Protoplanet both start the process with a cloud that will eventually spin. Also, somewhere in the process shrinking/compacting occurs. Howeve.

The continuous flow from the solar magnetic field and subatomic particles from the solar atmosphere, or corona, into the solar system, is called solar wind. .

In an atom a nucleus is surrounded by electrons that are in constant motion.The electrons are in electron clouds, or where they go in the atom. The proton, p.

In other words, ionized gas comprises plasma. This plasma can escape the sun’s gravity becoming the solar wind that is continuously blowing through the solar.


Why does hydrogen ionization happen in HII regions? - Astronomy

THE ATOMIC HYDROGEN EMISSION SPECTRUM

This page introduces the atomic hydrogen emission spectrum, showing how it arises from electron movements between energy levels within the atom. It also looks at how the spectrum can be used to find the ionisation energy of hydrogen.

What is an emission spectrum?

Observing hydrogen's emission spectrum

A hydrogen discharge tube is a slim tube containing hydrogen gas at low pressure with an electrode at each end. If you put a high voltage across this (say, 5000 volts), the tube lights up with a bright pink glow.

If the light is passed through a prism or diffraction grating, it is split into its various colours. What you would see is a small part of the hydrogen emission spectrum. Most of the spectrum is invisible to the eye because it is either in the infra-red or the ultra-violet.

The photograph shows part of a hydrogen discharge tube on the left, and the three most easily seen lines in the visible part of the spectrum on the right. (Ignore the "smearing" - particularly to the left of the red line. This is caused by flaws in the way the photograph was taken. See note below.)

Note: This photograph is by courtesy of Dr Rod Nave of the Department of Physics and Astronomy at Georgia State University, Atlanta. The photograph comes from notes about the hydrogen spectrum in his HyperPhysics pages on the University site. If you are interested in more than an introductory look at the subject, that is a good place to go.

Ideally the photo would show three clean spectral lines - dark blue, cyan and red. The red smearing which appears to the left of the red line, and other similar smearing (much more difficult to see) to the left of the other two lines probably comes, according to Dr Nave, from stray reflections in the set-up, or possibly from flaws in the diffraction grating. I have chosen to use this photograph anyway because a) I think it is a stunning image, and b) it is the only one I have ever come across which includes a hydrogen discharge tube and its spectrum in the same image.

Extending hydrogen's emission spectrum into the UV and IR

There is a lot more to the hydrogen spectrum than the three lines you can see with the naked eye. It is possible to detect patterns of lines in both the ultra-violet and infra-red regions of the spectrum as well.

These fall into a number of "series" of lines named after the person who discovered them. The diagram below shows three of these series, but there are others in the infra-red to the left of the Paschen series shown in the diagram.

The diagram is quite complicated, so we will look at it a bit at a time. Look first at the Lyman series on the right of the diagram - this is the most spread out one and easiest to see what is happening.

Note: The frequency scale is marked in PHz - that's petaHertz. You are familiar with prefixes like kilo (meaning a thousand or 10 3 times), and mega (meaning a million or 10 6 times). Peta means 10 15 times. So a value like 3 PHz means 3 x 10 15 Hz. If you are worried about "Hertz", it just means "cycles per second".

The Lyman series is a series of lines in the ultra-violet. Notice that the lines get closer and closer together as the frequency increases. Eventually, they get so close together that it becomes impossible to see them as anything other than a continuous spectrum. That's what the shaded bit on the right-hand end of the series suggests.

Then at one particular point, known as the series limit, the series stops.

If you now look at the Balmer series or the Paschen series, you will see that the pattern is just the same, but the series have become more compact. In the Balmer series, notice the position of the three visible lines from the photograph further up the page.

Complicating everything - frequency and wavelength

You will often find the hydrogen spectrum drawn using wavelengths of light rather than frequencies. Unfortunately, because of the mathematical relationship between the frequency of light and its wavelength, you get two completely different views of the spectrum if you plot it against frequency or against wavelength.

The relationship between frequency and wavelength

The mathematical relationship is:

Rearranging this gives equations for either wavelength or frequency.

What this means is that there is an inverse relationship between the two - a high frequency means a low wavelength and vice versa.

Note: You will sometimes find frequency given the much more obvious symbol, f.

Drawing the hydrogen spectrum in terms of wavelength

This is what the spectrum looks like if you plot it in terms of wavelength instead of frequency:

. . . and just to remind you what the spectrum in terms of frequency looks like:

Is this confusing? Well, I find it extremely confusing! So what do you do about it?

For the rest of this page I shall only look at the spectrum plotted against frequency, because it is much easier to relate it to what is happening in the atom. Be aware that the spectrum looks different depending on how it is plotted, but, other than that, ignore the wavelength version unless it is obvious that your examiners want it. If you try to learn both versions, you are only going to get them muddled up!

Note: Syllabuses probably won't be very helpful about this. You need to look at past papers and mark schemes.

If you are working towards a UK-based exam and don't have these things, you can find out how to get hold of them by going to the syllabuses page.

Explaining hydrogen's emission spectrum

The Balmer and Rydberg Equations

By an amazing bit of mathematical insight, in 1885 Balmer came up with a simple formula for predicting the wavelength of any of the lines in what we now know as the Balmer series. Three years later, Rydberg generalised this so that it was possible to work out the wavelengths of any of the lines in the hydrogen emission spectrum.

What Rydberg came up with was:

RH is a constant known as the Rydberg constant.

n1 and n2 are integers (whole numbers). n2 has to be greater than n1. In other words, if n1 is, say, 2 then n2 can be any whole number between 3 and infinity.

The various combinations of numbers that you can slot into this formula let you calculate the wavelength of any of the lines in the hydrogen emission spectrum - and there is close agreement between the wavelengths that you get using this formula and those found by analysing a real spectrum.

Note: If you come across a version of Balmer's original equation, it won't look like this. In Balmer's equation, n1 is always 2 - because that gives the wavelengths of the lines in the visible part of the spectrum which is what he was interested in. His original equation was also organised differently. The modern version shows more clearly what is going on.

You can also use a modified version of the Rydberg equation to calculate the frequency of each of the lines. You can work out this version from the previous equation and the formula relating wavelength and frequency further up the page.

Note: You may come across versions of the Rydberg equation where the n1 and n2 are the other way around, or they may even be swapped for letters like m and n. Whichever version you use, the bigger number must always be the one at the bottom of the right-hand term - the one you take away. If you get them the wrong way around, it is immediately obvious if you start to do a calculation, because you will end up with a negative answer!

The origin of the hydrogen emission spectrum

The lines in the hydrogen emission spectrum form regular patterns and can be represented by a (relatively) simple equation. Each line can be calculated from a combination of simple whole numbers.

Why does hydrogen emit light when it is excited by being exposed to a high voltage and what is the significance of those whole numbers?

When nothing is exciting it, hydrogen's electron is in the first energy level - the level closest to the nucleus. But if you supply energy to the atom, the electron gets excited into a higher energy level - or even removed from the atom altogether.

The high voltage in a discharge tube provides that energy. Hydrogen molecules are first broken up into hydrogen atoms (hence the atomic hydrogen emission spectrum) and electrons are then promoted into higher energy levels.

Suppose a particular electron was excited into the third energy level. This would tend to lose energy again by falling back down to a lower level. It could do this in two different ways.

It could fall all the way back down to the first level again, or it could fall back to the second level - and then, in a second jump, down to the first level.

Tying particular electron jumps to individual lines in the spectrum

If an electron falls from the 3-level to the 2-level, it has to lose an amount of energy exactly the same as the energy gap between those two levels. That energy which the electron loses comes out as light (where "light" includes UV and IR as well as visible).

Each frequency of light is associated with a particular energy by the equation:

The higher the frequency, the higher the energy of the light.

If an electron falls from the 3-level to the 2-level, red light is seen. This is the origin of the red line in the hydrogen spectrum. By measuring the frequency of the red light, you can work out its energy. That energy must be exactly the same as the energy gap between the 3-level and the 2-level in the hydrogen atom.

The last equation can therefore be re-written as a measure of the energy gap between two electron levels.

The greatest possible fall in energy will therefore produce the highest frequency line in the spectrum. The greatest fall will be from the infinity level to the 1-level. (The significance of the infinity level will be made clear later.)

The next few diagrams are in two parts - with the energy levels at the top and the spectrum at the bottom.

If an electron fell from the 6-level, the fall is a little bit less, and so the frequency will be a little bit lower. (Because of the scale of the diagram, it is impossible to draw in all the jumps involving all the levels between 7 and infinity!)

. . . and as you work your way through the other possible jumps to the 1-level, you have accounted for the whole of the Lyman series. The spacings between the lines in the spectrum reflect the way the spacings between the energy levels change.

If you do the same thing for jumps down to the 2-level, you end up with the lines in the Balmer series. These energy gaps are all much smaller than in the Lyman series, and so the frequencies produced are also much lower.

The Paschen series would be produced by jumps down to the 3-level, but the diagram is going to get very messy if I include those as well - not to mention all the other series with jumps down to the 4-level, the 5-level and so on.

The significance of the numbers in the Rydberg equation

n1 and n2 in the Rydberg equation are simply the energy levels at either end of the jump producing a particular line in the spectrum.

For example, in the Lyman series, n1 is always 1. Electrons are falling to the 1-level to produce lines in the Lyman series. For the Balmer series, n1 is always 2, because electrons are falling to the 2-level.

n2 is the level being jumped from. We have already mentioned that the red line is produced by electrons falling from the 3-level to the 2-level. In this case, then, n2 is equal to 3.

The significance of the infinity level

The infinity level represents the highest possible energy an electron can have as a part of a hydrogen atom. So what happens if the electron exceeds that energy by even the tiniest bit?

The electron is no longer a part of the atom. The infinity level represents the point at which ionisation of the atom occurs to form a positively charged ion.

Using the spectrum to find hydrogen's ionisation energy

When there is no additional energy supplied to it, hydrogen's electron is found at the 1-level. This is known as its ground state. If you supply enough energy to move the electron up to the infinity level, you have ionised the hydrogen.

The ionisation energy per electron is therefore a measure of the distance between the 1-level and the infinity level. If you look back at the last few diagrams, you will find that that particular energy jump produces the series limit of the Lyman series.

Note: Up to now we have been talking about the energy released when an electron falls from a higher to a lower level. Obviously if a certain amount of energy is released when an electron falls from the infinity level to the 1-level, that same amount will be needed to push the electron from the 1-level up to the infinity level.

If you can determine the frequency of the Lyman series limit, you can use it to calculate the energy needed to move the electron in one atom from the 1-level to the point of ionisation. From that, you can calculate the ionisation energy per mole of atoms.

The problem is that the frequency of a series limit is quite difficult to find accurately from a spectrum because the lines are so close together in that region that the spectrum looks continuous.

Finding the frequency of the series limit graphically

Here is a list of the frequencies of the seven most widely spaced lines in the Lyman series, together with the increase in frequency as you go from one to the next.

As the lines get closer together, obviously the increase in frequency gets less. At the series limit, the gap between the lines would be literally zero.

That means that if you were to plot the increases in frequency against the actual frequency, you could extrapolate (continue) the curve to the point at which the increase becomes zero. That would be the frequency of the series limit.

In fact you can actually plot two graphs from the data in the table above. The frequency difference is related to two frequencies. For example, the figure of 0.457 is found by taking 2.467 away from 2.924. So which of these two values should you plot the 0.457 against?

It doesn't matter, as long as you are always consistent - in other words, as long as you always plot the difference against either the higher or the lower figure. At the point you are interested in (where the difference becomes zero), the two frequency numbers are the same.

As you will see from the graph below, by plotting both of the possible curves on the same graph, it makes it easier to decide exactly how to extrapolate the curves. Because these are curves, they are much more difficult to extrapolate than if they were straight lines.

Both lines point to a series limit at about 3.28 x 10 15 Hz.

Note: Remember that 3.28 PHz is the same as 3.28 x 10 15 Hz. You can use the Rydberg equation to calculate the series limit of the Lyman series as a check on this figure: n1 = 1 for the Lyman series, and n2 = infinity for the series limit. 1/(infinity) 2 = zero. That gives a value for the frequency of 3.29 x 10 15 Hz - in other words the two values agree to within 0.3%.

So . . . now we can calculate the energy needed to remove a single electron from a hydrogen atom. Remember the equation from higher up the page:

We can work out the energy gap between the ground state and the point at which the electron leaves the atom by substituting the value we've got for frequency and looking up the value of Planck's constant from a data book.

That gives you the ionisation energy for a single atom. To find the normally quoted ionisation energy, we need to multiply this by the number of atoms in a mole of hydrogen atoms (the Avogadro constant) and then divide by 1000 to convert it into kilojoules.

Note: It would be wrong to quote this to more than 3 significant figures. The value for the frequency obtained from the graph is only to that accuracy.

This compares well with the normally quoted value for hydrogen's ionisation energy of 1312 kJ mol -1 .

Questions to test your understanding

If this is the first set of questions you have done, please read the introductory page before you start. You will need to use the BACK BUTTON on your browser to come back here afterwards.


The Birth of Disks Around Protostars

The dusty disks around young stars make the news regularly due to their appeal as the birthplace of early exoplanets. But how do disks like these first form and evolve around their newly born protostars? New observations from the Atacama Large Millimeter/submillimeter Array (ALMA) are helping us to better understand this process.

Formation from Collapse

Stars are born from the gravitational collapse of a dense cloud of molecular gas. Long before they start fusing hydrogen at their centers — when they are still just hot overdensities in the process of contracting — we call them protostars. These low-mass cores are hidden at the hearts of the clouds of molecular gas from which they are born.

Aerial image of the Atacama Large Millimeter/submillimeter Array. [EFE/Ariel Marinkovic]

But how do these Keplerian disks — which eventually have scales of hundreds of AU — first form and grow around protostars? We need observations of these disks in their early stages of formation to understand their birth and evolution — a challenging prospect, given the obscuring molecular gas that hides them at these stages. ALMA, however, is up to the task: it can peer through to the center of the gas clouds to see the emission from protostellar cores and their surroundings.

ALMA observations of the protostar Lupus 3 MMS. The molecular outflows from the protostar are shown in panel a. Panel b shows the continuum emission, which has a compact component that likely traces a disk surrounding the protostar. [Adapted from Yen et al. 2017]

New Disks Revealed?

In a recent publication led by Hsi-Wei Yen (Academia Sinica Institute of Astronomy and Astrophysics, Taiwan), a team of scientists presents results from ALMA’s observations of three very early-stage protostars: Lupus 3 MMS, IRAS 15398–3559, and IRAS 15398–2429. ALMA’s spectacular resolution allowed Yen and collaborators to infer the presence of a 100-AU Keplerian disk around Lupus 3 MMS, and signatures of infall on scales of <30 AU around the other two sources.

The authors construct models of the sources and show that the observations are consistent with the presence of disks around all three sources: a 100-AU disk around a 0.3 solar-mass protostar in the Lupus system, a 20-AU disk around a 0.01 solar-mass protostar in IRAS 15398–3559, and 6-AU disk around a 0.03 solar-mass protostar in IRAS 15398–2429.

By comparing their observations to those of other early-stage protostars, the authors conclude that in the earliest protostar stage, known as the Class 0 stage, the protostar’s disk grows rapidly in radius. As the protostar ages and enters the Class I stage, the disk growth stagnates, changing only very slowly after this.

These observations mark an important step in our ability to study the gas motions on such small scales at early stages of stellar birth. Additional future studies will hopefully allow us to continue to build this picture!

Citation

Hsi-Wei Yen et al 2017 ApJ 834 178. doi:10.3847/1538-4357/834/2/178


Data Collection

250 households in Tasmania were targeted for the Survey. Of the 247 households that ended up being used for analysis 122 were classed as rural/regional whereas 125 were located in &ldquoHobart&rdquo.

At this stage it should be noted that in order to &ldquoavoid a demographic bias in the sample&rdquo the following controls were applied to the National sample (

Firstly, these controls strike me as weird. Does it matter that half of the survey respondents are female? Does a child use a different amount of electricity to an adult, or &ldquoaverage person&rdquo? And really, do those with a higher education use a different amount of electricity? And even if it did make a difference, why not set the level at 24% multivariate linear regression analysis to model national electricity usage. This, honestly, sounds much more complicated than it is. Basically you just model overall electricity use as a collection of components:

In maths talk, this looks like:

Total Electricity Use = base level + (a x component 1) + (b x component 2) + (c x component 3) &hellip.

where the numbers a,b,c,&hellip. are obtained by the model

graphically, this looks like putting a collection of parameters into a &lsquoblack box&rsquo (the model), to obtain an answer as close as possible to the actual electricity use

Electricity use may well be described pretty well by just one or two parameters, or it may take five or six. In statistics terms, the model should improve (i.e. be able to model an energy use that&rsquos the same as actual useage) as more and more parameters are included.

For Tasmania, the top five variables (the parameters most significant in predicting electricity use), in order, are:

  • household size (number of people in house)
  • how often a dishwasher is used
  • whether halogen down lights are installed
  • whether the home has 1, 2 or 3 &ldquohot climate&rdquo appliances
  • how often and how long computers are used in the house

What we want our model to show how much of the behaviour of &lsquototal energy use&rsquo can be explained by each combination of input parameters.

If we just take the most significant parameter - the number of people who live in the house - the ACIL Tasman ** can explain 27% of the total fluctuation in overall energy use** between households in Tasmania.

If we include 14 parameters, their model can successfully estimate the bills for nearly half of Tasmanian households.


Ep. 255: Observing Hydrogen

Hydrogen is the most common element in the Universe, formed at the beginning of everything in the Big Bang. It’s the raw material of stars, gathering together through mutual gravity into vast nebulae. Astronomers can learn so much looking for hydrogen in the Universe. Here’s why, and how they do it.

Show Notes

  • Google+: Pamela and Fraser
  • Sponsor: 8th Light — Wiki — GSU — AstronomyKnowHow.com — Society for Popular Astronomy — Haystack Observatory — Smithsonian Astronomical Observatory — American Association of Amateur Astronomers

Transcript: Observing Hydrogen

Fraser: Welcome to AstronomyCast, our weekly facts-based journey through the Cosmos, where we help you understand not only what we know, but how we know what we know. My name is Fraser Cain I’m the publisher of Universe Today, and with me is Dr. Pamela Gay, a professor at Southern Illinois University-Edwardsville. Hi, Pamela. How are you doing?

Pamela: I’m doing well. How are you doing, Fraser?

Fraser: Doing really well … having fun recording another episode of AstronomyCast with all of our closest friends here on Google plus, so if you want to watch us live record the show, which we know not many people can actually do because they have jobs, and lives, and things like that, but yeah, you can just go to CosmoQuest.org/hang-outs and you’ll see a list of all of the shows that we do. We do a ton on astronomy-related content and science with us, and Phil Plate, Emily Lakdawalla from Planetary Society, and Alan Boyle from MSNBC, so we got lots of space friends and we’re doing a lot of really good content, so you should come and check it out, and that’s at CosmoQuest.org/hang-outs. We also…we embed the shows there so you can watch them live, you can participate in the conversations, and then, of course, if you can’t watch it live, we do try and mix everything and feed it into the AstronomyCast feed, and actually, I realized we’ve been putting the weekly space hang-out into the AstronomyCast feed and didn’t warn anybody, so…[laughing].

Pamela: [laughing] You suddenly have new content!

Fraser: Yeah! So if you’ve noticed now that you’re getting like an extra hour of audio content every week, that’s this weekly space hang-out that we’re doing on Google plus. No one’s complained, but no one has also said “Hey, thanks for putting that in there. I really appreciate that!” So I don’t know whether people are deleting them, or what. But if you’re getting those and you’re happy, that’s great if you’re getting them and you’re sad, then also let me know because we could also just break it up. You know, it’s pretty interesting, it’s the kind of content that people always asked us to do, but we never did, which is talk about the news and the current events and analysis of that kind of stuff, which is totally different from AstronomyCast, so anyway, that’s all in there. Sorry about that hope you’re OK with that. Please let us know if you’re not. Alright, well, why don’t we get cracking then?
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Fraser: So hydrogen is the most common element in the Universe, formed at the beginning of everything in the Big Bang. It’s the raw material of stars, gathering together through mutual gravity into vast nebulae. Astronomers can learn so much looking for hydrogen in the Universe. Well, here’s why and how they do it. Now, we wanted to, sort of, when we first sort of set up this show, I was like “OK, so the topic is hydrogen!” And you were like “No, no, no, that’s too big, that’s too much. Let’s just observe hydrogen.”

Pamela: It’s like 70% of the Universe. There’s a whole lot of stuff going on and…let’s keep focused.

Fraser: Like chemistry, and fusion, and powering cars, and things like that, so…but at least I think we should just have a brief conversation just about the formation of hydrogen and where it all came from, and then I promise we won’t go into the detailed chemistry of it.

Pamela: Well, so hydrogen — talking about its formation is somewhat silly. You take energy, you leave it on a shelf, it becomes protons probably (or other particles), and if it’s enough energy to become a proton, well, one proton that counts as ionized hydrogen, let it near a neutron, you now have a slightly more interesting hydrogen atom. Give it an electron — you now have a neutral hydrogen atom, so basically, hydrogen is that stuff that just formed when the Universe’s energy cooled off enough to start forming particles. Everything more complicated than hydrogen, you have to have some sort of nuclear fusion reaction take place in order to get to it, so hydrogen is just that simple thing that comes out of energy.

Fraser: And so back when the…during the Big Bang, when everything was just too hot, you just had raw energy…

Fraser: And then as things cooled down, that raw energy turned into protons, and…

Pamela: Protons, neutrons, electrons…

Fraser: …and electrons, and you just, you know, you just gather them together in the simplest possible way and that’s hydrogen. Obviously, we talked about it in a few episodes where you had this moment where the entire Universe was in this state of a star, and the hydrogen atoms were being fused into helium, and that’s where we get the helium from, but really, and then the expansion continued and now we’re just left with all this hydrogen, just this raw material, the building block of the entire Universe, so…and then why is it, then, I guess, important, then, for astronomers to observe hydrogen?

Pamela: Well, it’s not so much that it’s important to be able to observe hydrogen so much as we can’t help but observe hydrogen. It’s out there, and it’s causing us a whole variety of good things, and bad things, so on one hand, every time we’re looking at a star, we’re observing an excited hydrogen atmosphere. Every time we look at a beautiful nebula, we’re observing a cloud that’s rich in hydrogen gas that’s usually glowing red. When we start trying to look through the galaxy in radio light, we find all of the cold parts of space permeated with what’s called the 21-centermeter line of hydrogen. It’s just everywhere. Even when we look at high red-shift galaxies, we find in the spectra of these galaxies all of these places where intervening hydrogen gas has sucked the light out of the spectra of these distant galaxies, so if you study astronomy you’re just going to over and over come across the vocabulary of hydrogen. It can get a bit overwhelming, and that was actually part of the inspiration for this show. We’ve been doing live star parties, and I realized last night we’re talking , “H-II” — all of these different terms, and no one knows what the heck we’re talking about.

Fraser: Right, so hydrogen is the most abundant element in the Universe, so you just can’t help but see it everywhere you look.

Fraser: And so we might as well understand what it is that we’re looking at. Is it almost like all astronomers are pretty much hydrogen astronomers, you know? Like, a certain percentage of the time is just dealing with the hydrogen in everything they’re looking at?

Pamela: Yes, and one of the hazing rituals of getting a physics degree is learning all of the Quantum Mechanics of the hydrogen atom, and so by the time you finish getting even an undergraduate degree, you are intimately aware of the inner workings of hydrogen at levels you may not want, and you know how to find it all over the Universe.

Fraser: But you’re going to spare us the Quantum Mechanics today.

Pamela: I’m going to spare you the Quantum Mechanics today.

Fraser: OK – good, good. Alright, so then let’s talk about the different flavors of hydrogen that astronomers will observe out in the Universe.

Pamela: Well, the most common way that we confront hydrogen just as we peer through the sky with a pair of binoculars, or with a telescope is what’s called Hydrogen Balmer Lines, so when you look out, you’ll see particularly what’s called either Hydrogen Balmer Alpha Line, or just hydrogen-alpha because we get lazy. This is that bright red color that is associated with most nebulae, and it comes from the fact that the hydrogen’s energy levels are such that that one lone electron it’s got – it can jump between, well, it’s lowest energy level, to its second energy level, and transitions in and out of that lowest energy level. Those occur in the ultraviolet where we don’t see them with our eyes, so those are probably the most common transitions, but the ones we don’t see because ultraviolet gets blocked by our atmosphere. Now, go up one set of energy levels, and look at the transitions in and out of the second energy level. Well, there we have what’s called the 3 to 2 – from the third energy level to the second energy level transition — and that’s at this beautiful, red color that we see in “Open” signs at the local deli, and we see in all of these nebula that are all through the sky, so that red color associated with nebulosity – that is the lowest transition in and out of the second energy level of hydrogen, and this transition was discovered by a dude named Balmer, so it’s called the Balmer energy set, and alpha is for the lowest one, so 3 to 2 is alpha, then if you went 4 to 2 that would be beta, and so on through the list.

Fraser: And just to be clear, I think we talked about this in previous shows as well, right? This is that transition, that energy transition, right? When an atom of hydrogen, where it’s got its proton, it’s got its neutron, and then it’s got this electron, and that electron jumps up or down a level, you can get like a release of energy, and we’re seeing the photons streaming away from these nebula as these electrons are being released.

Pamela: So to get this to happen, you have to have a cloud of gas that’s getting heated up by something. So there’s either a bright star embedded in the cloud, there’s a whole bunch of bright stars embedded in the cloud, and the light from the stars is exciting the hydrogen so that it’s making this transition.

Fraser: Now, sorry, when you say just…I’m trying to be kind of precise here. So when you say exciting, you mean photons are streaming off of this star…

Pamela: Those photons are getting absorbed by the hydrogen atom.

Pamela: And the hydrogen atom in response to absorbing this photon, the electron is jumping to a higher energy level, and it might actually jump a whole bunch of energy levels, depending on what energy it gets hit with, and this actually has a neat effect where if the geometry is such that you look out, you look at the cloud and the star that you’re looking at is on the other side of the cloud, when you look at the cloud, you’ll actually see the hydrogen alpha light, that red light, removed from the colors that you’re looking at. Now, if instead, the star is off to the side and not precisely lined up, then you see that color that red energy from the star is getting absorbed by the hydrogen, re-radiated in all directions, and so you end up seeing the nebula as red.

Fraser: Right, but the point is (and this is where the whole concept of Quantum comes from, right?) that there is this very discreet, very specific step that these electrons take as they jump up the energy levels, and with it there is the corresponding release that comes out in a very specific color, and it’s that color of radiation that we see with our telescopes, and that astronomers are really specifically looking for. They’re actually…they’re limiting the entire spectrum that they could see down to that exact, specific light.

Pamela: And this is actually something that anyone out there listening can experience for themselves. A lot of gag stores, a lot of novelty stores will sell these prism glasses that create rainbows when you look through them. Well, if you get one of these pairs of rainbow glasses, and you walk up to your local deli, you walk up to your local pub, whatever, and you look through these glasses at the neon signs, you’ll see the discreet, specific lines given off by the atoms in that sign, so if you look at a red “Open” sign, you’re going to see this bright red line that comes from the hydrogen alpha, but you’ll also see this gap, and then this bright (they call it “cyan,” to me, I’d call it turquoise)…this bright turquoise line, and that’s hydrogen beta. Then a little bit over to the side from that is hydrogen gamma – this is the 5 to 2 transition (and this is like Crayola blue, or that 00255 if you work in RGB colors), and so you’ll then start seeing closer and closer-spaced, deeper shades of blue as you look at the spectra of that red “Open” sign, and then you’ll see a completely different set of fingerprints if you look at a green sign, or a purple sign, but that red “Open” sign has this distinctive spectra through the novelty rainbow glasses that’s the Hydrogen Balmer series.

Fraser: Right, so I guess what astronomers are doing, right, is they’re filtering out every color of light except for that specific, sort of, in the frequency range that they’re trying to see. The equivalent of putting those crazy glasses on…

Pamela: If we use a hydrogen-alpha filter, yeah.

Fraser: Right, and so that’s the point, right? Astronomers will have a collection of these filters. They’ll have one for hydrogen alpha…how many hydrogen-related filters will astronomers use?

Pamela: So at a certain point, you stop using filters and you start doing imaging spectroscopy, so it’s not too uncommon to have a H alpha filter, a Lyman-alpha filter if you’re working in the ultraviolet, or what will also happen is since these lines are given off by galaxies at different red-shifts, people will actually create special filters tuned to only detect, say, Lyman-alpha. This is the 1 to 2 transition in hydrogen that if it’s nearby we can’t see because it’s UV, but if a galaxy is far away, and its light is getting shifted into the red, that color that’s usually so blue we can’t see it – it gets moved a little bit redder, a little bit redder, a little bit redder until we can see it, and they’ll create filters tuned to see the Lyman-alpha of galaxies that are moving at specific velocities.

Fraser: And I guess this is part of the thing where the amount of that frequency is so tight that if it is red-shifted, you’ve got to push it up and down the frequency. So astronomers know that they want to see this specific kind of frequency of light, and they’ve got the tools to be able to see it, but what does seeing it tell them? Why do they want to do this?

Pamela: Well, it’s… it depends on what you’re doing.

Fraser: Trying to do science.

Pamela: [laughing] And so the thing is there’s lots of different science that you could be doing. For instance, when we’re looking at different nebulae locally, we’re often trying to figure out what is the distribution of temperature in a cloud of gas, what is the density of the gas, and so when we’re looking at the hydrogen alpha light, when we’re looking at the light in all of these different energy levels of hydrogen, what we’re trying to do is figure out just how hot is that gas. And this is where we start talking about things like H-II regions. So an H-II region…the crazy notation we use in astronomy is a letter from the periodic table is clearly the abbreviation for the atom, if it has a Roman numeral “I” next to, that’s something that hasn’t been ionized at all — it’s completely neutral. If it has a “II” next to it, that means we’ve yanked off one electron. If it has a “III” next to it, we’ve yanked off two electrons. So take the number, subtract one, and that’s how many electrons we’ve removed from the atom. So when we’re talking about the H-II region, we’re talking about a region of space filled with hydrogen gas, and that gas is ionized one time to remove that one electron. Now in these H-II regions, this is a cloud of gas that is typically being heated up by really hot, bright stars, so when you look at the Orion nebula with all of it’s O-giant stars embedded in the gas, you’re looking at an H-II region, and in these regions the hydrogen atoms will periodically glom on to one of these free electrons, and as they glom on to the free electron, the electron will cascade down through the different energy levels, and it will give off hydrogen alpha, it will give off hydrogen beta, it will give off all these different parts of the spectrum, and by looking at that, and looking at the ratios of how many of the atoms appear in the different energy levels, we can start to get at the density of the material and the temperature of the material.

Fraser: Now, you mentioned a couple of other things as well. And there’s neutral hydrogen, and cold hydrogen, and those are useful for astronomers to observe as well, right?

Pamela: Right, and so another one of the things that we look at is what’s called the 21-centimeter line of hydrogen, and this is perhaps one of the harder things to try and explain. It’s actually something that when we teach it, we talk about this is something that was originally referred to as “Not going to happen, never going to be observable…” and it’s because it’s a process that takes a long, long time for it to happen, so if you take a hydrogen atom, its proton in the center has what we call in Quantum Mechanics a “spin,” and the spin is either spin up or spin down, and its orbiting electron has the same thing. It either has a spin up or a spin down, and ideally the two little bits — they want to be lined up the same, and so what you’ll have is if you leave hydrogen alone long enough, and it’s not in its lowest possible energy, you’ll end up getting that “spin-flip” and the energy given off in this flip is energy that corresponds to light with a wavelength that’s 21 centimeters long. Now, the probability, in most cases, is that before the atom has a chance for that flip to take place (because it takes a long time for the atom to finally get around to flipping probabilistically), it’s probably going to undergo a collision, it’s probably going to undergo and excitation – something’s going to happen to it. The only way that you’re going to consistently get this spin-flip is if you have a whole bunch of gas, it’s really cold, and thus not moving, so all the little atoms are just sort of going, “not moving, moving very slowly…” and it’s very diffuse gas as well, so you need cold, diffuse gas.

Fraser: Well, that’s kind of interesting though, right, because there’s a way…like, you wouldn’t think if it’s out there, just super-cold in space, just sitting there, not interacting, you would think there’d be nowhere to see it, it would just be invisible, but because there’s this crazy Quantum effect, they just randomly spin-flip, you get a release of radiation that’s very subtle, but it’s there and let’s you detect it.

Pamela: And so this is one of the ways we’re able to measure the rotation rate of our galaxy out to extremely high radii. So what we do is we use radio telescopes, and this is actually the type of thing that undergrads can do, or any amateur who builds their own at-home radio dish, and you can get kits to do that. This is an experiment you can do is identify where the clouds of cold gas are out in the outer wings of the arms of the Milky Way, take a look at them, and measure the Doppler shifting of that 21-centimeter line, and from the Doppler shifting you can get the rate at which the cloud is moving forward and backward in that direction in the sky, and you can use geometry then to start to then get at the orbital velocity of this gas and at the end of the day, this gives you the rotation curve for our galaxy that shows that everything is moving at about the same velocity as you move out toward the outer parts of the galaxy, and thus, you can demonstrate for yourself there is something gravitationally changing. This is dark matter.

Fraser: Well, I think that should be everyone’s homework for this week, then. So everyone should go out and observe the 21-centimeter line, and calculate the Doppler shifting, use geometry to determine the motion, the rotational motion of our position within the Milky Way.

Pamela: Completely elementary!

Fraser: Completely elementary – everyone, get on that! So what are these cold…? I mean, OK, so we can use these cold clouds of gas as weigh points, as places to determine position, but I mean, aren’t these future nurseries of stars?

Pamela: Not necessarily. The thing is that in order to get a star-forming region, you have to have dense gas that has sufficient mass that when you collapse it down and things start forming, you get enough mass leftover to form a star, and some clouds of gas just aren’t massive enough that they’re ever going to form anything meaningful, and in other cases, the clouds of gas as they are right now are so diffuse and so stable that we don’t see star formation in their immediate future. Now, spiral arms do help trigger star formation because what ends up happening is as these clouds of material orbit around the Milky Way, they get pulled in on the one side to the spiral arm, and then as they try and orbit out the other side of the spiral arm, they get slowed down, and as they linger in the spiral arm, there’s a good chance that there’s going to be collisions, there’s going to be compressions, there’s going to be shock waves from supernovae, and all of these effects may cause some of these otherwise far-too-diffuse clouds of gas to have star formation, but in general, our galaxy’s only about 1% effective at transforming gas into stars.

Fraser: So astronomers don’t see…like, don’t really do a lot of searching for great, big clouds of future nurseries. It’s more like waiting until the…you know, I guess it moves into to that hydrogen alpha phase, where you’re actually starting to see the light coming off the nebula that you start to identify these star-forming regions?

Pamela: Well, there’s lots of things that we do look at, and we’re like, “THAT is forming stars right now,” and this is where people who work in the radio and the millimeter, they actually start mapping out some of these clouds. So there are certain, what are called “bok globules.” These are extremely dense, often molecular hydrogen regions, so this is the other form of “H-two” that when you’re doing an audio show, it makes no sense. So you have “H-Roman numeral II,” which is ionized hydrogen, and you have “H-subscript 2,” which is molecular hydrogen, and when you look at these dense, black regions on the sky (Horsehead Nebula isn’t a bok globule, but it’s an example of one of these dense, black regions on the sky)…when you look at these dense, black regions in the sky in the optical, they just look like the never-ending story, “Great Nothing,” ate a part of the Universe, but when you start to look at them instead in millimeter wavelengths, you start to see they’re knots of thermally-radiating areas. These are areas where the gas has begun to contract, and as the gas squishes down, the atoms start hitting each other and this process radiates away, basically, warmth. So this is infrared this is millimeter to light. You can sort of think of this as if you rub your hands together, it’s going to generate heat, and if you had an infrared camera, you could actually hold your hands up and see that change in temperature from rubbing your hands. Now, when the gas starts colliding like that, you start initially giving off in the radio light. Now, you wait as it continues to collapse, stars start to form, starts to light up in the infrared, and eventually it brings itself all the way into the bright blue UV when you get the youngest stars actually igniting, but so we look for those dark, molecular clouds that are high-density, and those…yeah, they do probe those for star formation, but not every blob of gas is necessarily going to form stars.

Fraser: Can we look for places where, like, hydrogen is absorbing light? You know, like we look for places where certain elements are actually blocking, right?

Pamela: And so when we look at nebula, we talk about there being reflection nebula, we talk about there being emission nebula, and the truth is it’s just a matter of geometry. So if it’s star-cloud-observer, that cloud is going to absorb out the hydrogen lines. If it’s cloud in front of us, star off to the side, then we see emission lines, and so there’s lots of different ways, and it’s all about geometry that controls what we’re able to see.

Fraser: And I think as we’ve been really experiencing with doing these live star parties, and we have one person, we have Gary, who has got this just phenomenal 14-inch telescope, but he’s in this really polluted area — he’s in Los Angeles — and yet he seems to be able to pull together these really sensitive images of nebula. So, why does this hydrogen look so crisp and clear even when you’ve got really bad polluted skies?

Pamela: So he’s cheating in a way. If you’ve ever had one of those kids’ toys, or cereal boxes where you get the little red filter, and you look at this scrambled mess on the side of the cereal box, and then when you put the red filter in front of it you suddenly see a message. Well, what’s happening is, in that case, is you have all this visual noise, and that visual noise gets removed when you put the red filter in front of it — and Gary’s doing the exact same thing. In his case, he’s in the Los Angeles basin, and there’s for the most part sodium lights (those are the yellow parking lot lights that make the sky glow this raspberry color on a cloudy night), and then there’s also…now we’re getting more and more fluorescent lights which are giving off their blue UV light, and all of this is scattering skyward. Sometime it’s because they’re using stupid light fixtures that point the light upwards, or they’re illuminating buildings and it points the light upwards.

Fraser: Hate those people…

Pamela: Right, and sometimes it’s just a matter that you’re shining light down on cement, and the cement reflects the light back up however, the light’s getting upwards, it’s primarily consisting of the sodium light from the sodium light fixtures, and white light that’s peaking off in the UV from the fluorescents or peaking off towards the UV, not actually in the UV, and what he’s doing is he’s saying, “OK, I’m going to look at the sky, and I know that most of the sky is being brightly lit up by the atmosphere reflecting the sodium, and all of this white stuff that is peaking towards the blue. I’m going to try and get rid of as much of that as possible, and I’m going to focus in on one line of light – the hydrogen alpha light that’s in the red, as opposed to the blue, and the sodium’s yellow…” and by focusing on just that one color, well, suddenly, his background goes to black again because these street lights aren’t giving off hardly anything at all in , so suddenly the light pollution for the most part has been filtered out the same way all that visual noise was filtered out on the cereal box, and what’s left behind is only the hydrogen alpha light. Now, the crazy thing is if he actually went to a dark site, he’d get even more amazing images if he was able to use broader-band filters that were letting in more light all at once, but he does what he can, and he’s found a way to do really good astro-photography in a very light-polluted part of north America.

Fraser: Yeah. So there’s hope for all of us.

Pamela: Yeah, there is.

Fraser: Cool. Well, I think that about wraps it up for this week, so thanks a lot, and we’ll talk to you next week.

Pamela: That sounds great. Talk to you later, Fraser.

This transcript is not an exact match to the audio file. It has been edited for clarity.