# Amateur Astronomy

#378

Finally caught this. Interesting. Pity they had to counterbalance it with Lawrence Krauss being a complete oaf (as usual).

#379

The Strange Case of the Missing Dwarf

A little addendum to the exoplanet discussion (above). I mentioned that direct imaging of exoplanets was already under way. A new instrument called SPHERE was added to the European Southern Observatory’s VLT telescope on Cerra Paranal in Chile last summer. It’s an infrared detector which uses a coronagraph to block out the glare of the primary star and looks for the dust disk and planets surrounding it.

SPHERE (a) under test at Grenoble, (b) being hoisted into position at Paranal (anyone else think it looks a bit like a grand piano?), © installed at the VLT focus (large blue adapter ring); SPHERE is actually a suite of instruments mostly housed inside the black cover … those attached portable liquid nitrogen dewars (and spares at top right) for cooling the sensitive instruments are more common at the VLT than beer kegs in a pub!

SPHERE is actually a suite of instruments: a visible light imager, and an infrared spectrograph among other things. It has a number of tricks up its sleeve to make sure it gets the sharpest pictures of circumstellar regions. First, because the field of view is so small it has its own very precise adaptive optics instead of relying on the VLT’s – it can shift its own lens arrangement using servo motors at millisecond intervals to compensate for atmospheric distortion. The coronograph and associated Lyot stop make sure that stellar light is strongly rejected. And it uses polarimetry to further distinguish between direct stellar light and reflected light from a planet which has a different polarisation, very much like the clearer view you get with polaroid sunglasses.

So it’s interesting to note that SPHERE’s first attempt to cut its teeth on an unseen stellar companion has failed to come up trumps. The first science paper based on SPHERE observations is about the binary star V471 Tauri. The behaviour of the stellar orbits implied another object disturbing them, expected to be a “brown dwarf” – a type of failed star that didn’t quite ignite. It should be easily big and warm enough to be seen by SPHERE but nothing was observed. It’ll be interesting to see if SPHERE confirms or denies other expected objects that have previously failed to show up in the infrared, such as the disputed exoplanet, Fomalhaut b.

#380

Anyone see the beautiful not-quite-new “fingernail” crescent moon after sunset yesterday, near a brilliantly blazing Venus (beside which you could also see Mars if you looked carefully)?

The moon’s orbit is at that place in its 19 year cycle where it is ascending through the ecliptic just as our vernal equinox takes the sun over the equator. The orbit of the moon varies about five degrees above and below the ecliptic plane – the plane of the earth’s orbit around the sun. Due to axial tilt that’s quite different to the earth’s equatorial plane much of the time, except near the equinoxes.

The result of the current alignment is that for the next while, the sun and moon are over the equator, pulling on the earth in precisely the same (or opposite) directions whenever we have a new moon or full moon. (The moon is on the same side of us as the sun when new, and opposite it when full, but both arrangements have similar tidal effects). We are also having a series of “supermoons” when the moon is at perigee, its closest approach to the earth. Today it is just a hair over 220,000 miles away, which will be surpassed only by a supermoon in September when it will be 200 miles closer.

The combination of all these factors means that we’re about to have our highest tides for twenty years. Fortunately it doesn’t look like we’ll have adverse weather conditions – the combination of high tides, high rivers, onshore winds, and a low pressure system producing a large ocean swell is the “perfect storm” for causing extreme coastal flooding. The Dublin floods of 2002 produced huge damage with tides lower than those of this weekend. Yesterday (Saturday) the 15.2 ft tide was only a few inches lower than the maximum recorded range for Dublin.

#381

Seen it
Yeah it looked fantastic especially when lowest on the horizon
Weather conditions look nasty enough. May cause some flooding especially in the likes of Ennis (tidal) and the west coast

#382

I noticed that – the moon illusion seemed particular powerful with the crescent.

#383

Wannabe Stars Must Be Cool and Pass the Jeans Test

At the denouement of the Bond movie Goldfinger, the villain foolishly fires a gun inside an aeroplane, blowing out a window with predictable results, while Bond is saved and gets the girl. So one of them gets their just deserts and the other is sucked to oblivion.

Ok, the movie scene is overwrought and unrealistic, but it plays on the true fact that air will flow from a region of higher pressure to one of lower pressure. But if that’s the case, why doesn’t the high pressure air at sea level rush upward to where the air is thinner? And why doesn’t the air at the highest level get sucked away by the vacuum of space?

You probably know the answers. The air at sea level is pinned down by the weight of the air above it. That’s why the air is thinner higher up – there’s less air above it and therefore less weight on it. If you stood at the top of Mount Everest, two thirds of the earth’s atmosphere would be below you. The troposphere is the lowest layer of the atmosphere, in which all our weather occurs. At the top of the troposphere, a little higher than the operating ceiling for most commercial planes, nine tenths of the atmosphere is underneath.

And yet, the atmosphere is far from finished at that level. You have to go ten times higher, to 100 km, before you are officially “in space”. And even there, you are not in a hard vacuum. Satellites in low earth orbit, such as the International Space Station and Hubble Space Telescope at 400 km and 550 km respectively, have to be boosted higher at regular intervals or they will be eventually dragged down by the wisps of earth’s atmosphere.

So how about that other question about why the space vacuum doesn’t suck the air away? The fact, of course, is that vacuums don’t “suck”. Rather, air pressure pushes. At the microscopic level, air molecules are in constant motion due to heat energy, and air pressure is caused by collisions of molecules with each other and the walls of any container. The Maxwell-Boltzmann distribution tells us how the energy is shared, and yields the surprising fact that the average air molecule at room temperature is doing nearly a thousand mph between collisions! That’s fast … but not nearly fast enough to escape earth’s gravity. Several hundred kilometres up, gravity is not much different to what it is at ground level … and the velocity needed to escape earth is still well above 20,000 mph.

Air, like any gas, is highly compressible. The air at sea level is compressed by a kilogram of air above every square centimetre of it, the weight of a column of air stretching to the top of the atmosphere. We find that a mass of air transported from sea level to the top of Mount Everest occupies a larger volume. For a more practical experiment observe the foil top of an unopened pack of Ryanair Pringles bulge upward in an aeroplane at cruising altitude ( – the plane’s cabin being only pressurised to about 2000 metres altitude, 80% of the pressure at sea level where your Pringles were packed).

Due to the accumulation of weight with depth in the atmosphere, and the compressibility of the air, it is not surprising to find a logarithmic relationship between the air pressure and altitude. Note the linear height scale on the left, but pressure going in powers of ten on the right:

This is the case for any atmosphere supporting its own weight in so-called hydrostatic equilibrium where it neither compresses nor expands on average. Ok, here comes the maths bit. We find the relationship:

$\frac{\text{d}p}{\text{d}z}=-\rho g$

… which says that the rate of change of pressure with height is related to the density of the air times the gravitational acceleration. (Note: that’s a p for pressure on the left, but a Greek letter rho for density on the right). Of course, the air density itself is pressure dependent, so with a bit more maths we end up with:

$p\propto e^{-z/H}$

This is saying that the air pressure is inversely proportional to e (the base of natural logs, hence our logarithmic dependency) raised to the power of the height (z) divided by a value which is constant for a given atmosphere, and is referred to as the scale height of the atmosphere – the height over which the pressure changes by a factor of e. The scale height constant incorporates several others, including the strength of gravity, the temperature, and the molar mass of the particular gases involved. Lighter gases would result in an increased scale height, as would lower gravity or higher temperature.

You may be wondering at this stage what any of this has to do with astronomy, let alone the cryptic article title. Well, let’s cut to the chase. Suppose we were able to dial down the mass of the earth. What would happen to the atmosphere? Gravity would be weaker so the scale height would increase and the atmosphere would expand. Suppose we dialled the mass of the earth all the way down to zero. The atmosphere would just float away, right? Well, maybe not. The atmosphere has a mass of its own, and therefore generates its own extremely weak gravity. Would that be enough to hold the atmosphere together, albeit greatly expanded? That depends on the other crucial variable – temperature. Remember that this is what determines the speed of the gas molecules, and ultimately whether they reach the necessary speed to escape the clutches of gravity. The atmosphere’s own gravity is extremely low, so if we want it to hang together we’re going to have to make it incredibly cold.

And this brings us to how stars form from giant clouds of gas. Now we are talking about an ultra-extreme version of the scenario we just discussed. The cloud may be many light years wide. It is incredibly tenuous – thousands of times more rarefied than even the few molecules of atmosphere that drag on the space station. Nevertheless, such a vast volume may contain an appreciable mass of gas. Newton’s gravitational law tells us that the acceleration due to gravity at a given radius within a spherical gas cloud is determined only by the mass enclosed inside that radius. The influence of all the gas outside magically cancels out. Where G is the universal gravitational constant, the acceleration due to a mass M of radius r is:

$g=-\frac{GM}{r^2}$

… and we can take this value for g, and using r instead of z for height we can write the condition for hydrostatic equilibrium just as we did for earth’s atmosphere:

$\frac{\text{d}p}{\text{d}r}=-\frac{\rho GM}{r^2}$

Once again we are saying that the rate of change of pressure with radius is equal to the density times the gravitational acceleration. Once again, the density is affected by temperature and pressure. For an interstellar gas cloud in hydrostatic equilibrium, something must change if it is to collapse under its own weight to form a star. One factor is temperature. The gas cloud can radiate away heat energy to upset the balance by reducing its internal pressure. It can take tens or hundreds of millions of years for this to happen, depending on the efficiency of heat loss.

One interesting fact here concerns the chemical make up of the cloud. I mentioned in a previous article on spectroscopy that incandescent materials like metals can radiate at a large range of wavelengths, whereas an excited gas will only radiate at one or a small number of characteristic wavelengths. This makes metals much more efficient at losing heat. In the early universe the only material around was hydrogen and helium made in the Big Bang and it took a long time for the first generation of stars to get going. But the first generation of stars made heavier elements in their cores, and stellar winds and supernova explosions scattered this material, contaminating the interstellar medium with it. That means that subsequent generations of stars can form more quickly because modern gas clouds containing metals can cool quicker.

Other triggers for gas cloud collapse can be density changes brought about by shock waves from supernovae or disturbances from collisions between galaxies. We can see “starburst” galaxies containing huge numbers of young blue stars, evidence of recent galaxy mergers or collisions.

On the other hand, raising the temperature can cause a gas cloud to disperse. Ironically, this happens just after stars have formed. The radiation from the new star heats the surrounding medium, so the remnants of the cloud that gave birth to it are blown away. Often in regions of new star formation we see some stars still inside the “stellar nursery” perhaps visible only by their infrared light which can penetrate the gas, while other stars have drifted out and broken free, or have blown a big empty bubble around themselves as their stellar winds drive away gas.

If you know how to locate the constellation of Orion which is very strikingly visible to the south early in our February evenings, you’ll know that the hunter has a conspicuous belt of three stars, from which hangs a sword of three more stars. On closer inspection (but still visible to the naked eye) you can see that the middle “star” of the sword is strangely fuzzy, and it turns out to be this monster when viewed in a large telescope:

This is a region of active star formation, similar to where our own sun was born five billion years ago in the same spiral arm of the galaxy. The Orion nebula is something over twenty light years across and contains about 2,000 young stars. Stars with common origins in giant gas clouds form open clusters, which break up and drift apart over millions of years. The Pleiades constellation is another such open cluster. Even the Plough / Big Dipper consists mostly of stars that formed together about half a billion years ago. (Most constellations, of course, consist of stars that just happen to be in the same line of sight, but may not be associated with each other at all). Eventually, stars like the sun end up drifting alone in the galaxy, showing few signs of their nebular origins, except that the planets and other debris of our solar system are made up of the contaminants that were mixed with the original gas cloud and helped it cool.

The British astronomer Sir James Jeans, working from the equation for hydrostatic equilibrium, figured out the mass that a gas cloud has to have in order for it to collapse. This Jeans mass criterion, of course, depends on the size of the cloud and crucially on its temperature. Most gas clouds have to get down to an incredibly cold -250 °C or less before collapse occurs. Hence our title: if you wannabe a star you have to stay cool and pass the Jeans test. A great deal more could be said about star formation, but that’s quite long enough for one article.

#384

Surely when stars first formed the temperature was into the thousands kelvin?

#385

Good question, Carrot, to which the answer is: yes and no, depending on what you mean. If you mean that the temperature in the early universe was high, then no. By one billion years after the Big Bang the temperature has already fallen to 19 K. We’re not exactly sure when the first stars formed, but the latest information (just this month, and reported just the other day on what must’ve been a slow news day for the BBC) based on Planck telescope observations suggests the first stars formed after about 550 million years. The universe is slightly warmer then, but still a very frigid 28 K (-245 °C).

But yes you are right that the first stars must have formed from gas clouds at much higher temperatures than most of today’s stars. When a gas cloud starts to collapse, it only collapses until hydrostatic equilibrium is restored. That happens because the temperature starts to rise as the gas is compressed, according to the ordinary ideal gas laws. We can only get further collapse by radiating away more energy to cool the cloud again. I mentioned that the primordial material is particularly inefficient at cooling. The problem we then get is that the cooling time is way, way too long for us to have gotten stars as early as we did.

I should clarify that the Jeans mass is not the mass at which a cloud starts to collapse, perhaps to be halted again when the temperature rises, but the mass at which the cloud is guaranteed to collapse all the way down to a star. It is calculated by looking at the total thermal energy of the cloud, and comparing it to the (negative) gravitational potential energy. If these sum to less than zero (i.e. the gravitational potential is greater than the thermal energy), the cloud is gravitationally bound all the way down to star formation, at which point hydrostatic equilibrium is restored by the ignition of core fusion.

So the solution is that the gas clouds must have been much bigger for those early stars than is the case today. The higher gravity means that the cloud can exceed the Jeans mass at a higher temperature. The first stars must have been more massive than our sun by factors of hundreds, perhaps even a thousand. (There’s an upper limit on the mass of a star somewhere up around that level, because the radiation pressure from fusion in the core is so high that it lifts the entire envelope off the star and completely disrupts it – we think [*Wolf-Rayet stars * (https://en.wikipedia.org/wiki/Wolf–Rayet_star) may be such naked cores. On the other hand, giant clouds don’t necessarily form single stars because of possible fragmentation during collapse, but that’s a whole other story).

Today we can still get large stars forming from gas clouds at higher temperatures. But the gas clouds are able to cool much more efficiently, and the huge majority of stars are of low mass, smaller than our sun. So most stars (but not all) start out from very cold clouds.

I didn’t give the formula for the Jeans mass earlier. Here it is:

https://latex.codecogs.com/gif.latex?M_J%3D\frac{3kT}{2G\overline{m}}R

… where k and G are constants (the Boltzmann and gravitational constants), T is the temperature, R is the cloud radius and https://latex.codecogs.com/gif.latex?\overline{m} is the average mass per particle (which depends on the type of gas involved). The Jeans mass is the minimum mass above which a cloud must inevitably collapse. Because the formula also involves the radius of the cloud we can rearrange it to give a Jeans density (which I won’t bother showing). Exceeding the Jeans density also causes collapse, and I mentioned that overdense regions can be created by various dynamical processes.

Here’s an interesting diagram of temperature plotted against density, which gives you an idea of different cloud temperatures that could be accommodated for clouds of different masses at a range of initial cloud densities. As you can see, the range of cloud densities varies by a factor of about a million. The Jeans mass lines shown assume the cloud composition is neutral hydrogen. For very massive protostars (100 solar masses) at the highest densities, you could have cloud temperatures up to what looks like around a balmy 170 K (-100 °C). I don’t know if the low mass end of the lowest cloud densities can physically form stars, since the indicated temperature is less than a tenth of a Kelvin … colder than the CMB!

As for the very high mass clouds which might be relevant to the early universe, the diagram doesn’t tell us but if we assume the 1,000 solar mass line sits equally spaced above the other lines, it looks like you could extrapolate the temperature to maybe a thousand K at the high density end. That would make your supposition plausible, but I honestly don’t know. However, Google being our friend, I found this Sci Am article which confirms a temperature of several hundred to a thousand K for the primordial gas clouds (depending on cooling mechanisms). But bear in mind it’s based on modelling from 2002, and a lot changes in cosmology in more than a decade, as the Planck observations show.

#386

Cool. Thanks PS. I was looking up some more stuff on star formation and they are cathegories into population I, II and III stars. Population II stars appeared around 8.5 to 9 billion yrs ago. Ours is a population I system. While the search for solar systems is only in its infancy, are there any planetory systems associated with population II stars?

#387

You’re poking at the limits of my knowledge there (and I’m no better at Googling than the next man ). But as a general observation I’d say you don’t have to scratch too deep under the surface of most topics in astronomy before you find a lot of “known unknowns”. This is probably one of them.

It’s a fair approximation to say that all of exoplanets for which we have any detailed information are the result of the Kepler mission. And most of those are within in a few dozen light years, and practically all of them within a thousand light years. That puts them all, relatively speaking, in our back yard. Here’s the size and direction of the Kepler search space relative to our galaxy, but in reality a high proportion of the discovered planets would fit inside those crosshairs centred on the sun:

That puts all the discoveries inside our own Orion spiral arm of the galaxy, or pretty close by, so they’re all going to be population I stars. So I’m not sure the question of whether population II stars have planets can be answered yet. The conventional wisdom is that the more metal-rich Pop I stars are more likely to have planets for obvious reasons, the classification into populations I and II being defined by metallicity.

However, there seems to be a great deal yet to learn about the range of metallicities found in stars. Some are anomalously metal poor, even among the pop I stars, and one astronomer I met was devoting all his research into anomalous lithium deficiencies in just one stellar type. There’s also debate about how the early galaxies became metal enriched. If the theory I mentioned of the early population III stars being oversized is correct, then there should have been a lot of Type II (core collapse) supernovae, but not a lot of type Ia (old stars feeding on binary companions). That would effect the type of metallicity in the subsequent type II populations, not just the amount. Metallicity is usually characterised by just a single number, z, which is the proportion of atoms heavier than helium. (To an astronomer, everything heavier than helium is a metal). Pop II stars show broadly the right type of metallicity (oxygen, neon, silicon, magnesium) corresponding to having been enriched by the hypothesised earlier generations of stars.

My own brief study of quasar spectra showed plenty of metal lines. (Quasars are the luminous central regions of active galaxies, presumed to be the radiation from accretion discs around supermassive black holes). I’m not sure how that squares with the idea that supermassive black holes formed early in galaxy evolution. Studies of quasar metallicities have thrown up a bunch of paradoxes. Many quasars are far more metal rich than our sun, with z = 3 x solar. Quasar metallicity is proportional to luminosity, from which it is inferred that the presence of supermassive black holes causes high rates of star formation, the bigger the black hole the more stars, and it is the presence of these stars that causes the metal enrichment of the quasar host galaxy and the quasar itself.

But, then it turns out that when we look at high red shift quasars, once we allow for the luminosity-metallicity relation there doesn’t seem to be any evolution of metallicity with red shift (and thus age). It looks like the high metal content was there very early on. Studies show no change in metallicity of quasars out to red shift 6.4, which is a lookback time of nearly 13 billion years, to when the universe was less than a billion years old. Where did the metals come from?

Here’s an example quasar I looked at from the Sloan Digital Sky Survey. This one’s from a “mere” eleven billion years ago (red shift 2.35) and you can see prominent lines of carbon (as well as silicon, oxygen, magnesium etc.). Carbon cores are characteristic of intermediate mass, relatively long lived stars, so it’s expected to take about a billion years to enrich a galaxy with carbon. We get away with it for this example, but we also see these carbon lines in quasars at red shift 6, when the universe was less than a billion years old. That doesn’t make sense. I think we have a lot yet to learn on this topic.

#388

I mentioned that the moon and sun are pulling together and raising high tides, owing to the moon’s orbit this year taking it especially near the equator in the same season as the sun’s equinocteal passage across that line. This celestial coincidence also increases the chance of eclipses of both sun and moon, as the earth sits in a straight line with both of them.

Thus, on March 20th, just twelve hours before the vernal equinox, a total solar eclipse will cross the north Atlantic, passing through the Norwegian archipelago of Svalbard and moving south west across the Faroe Islands. Although we won’t quite get the totality in the British Isles, the northern tip of Scotland will see the sun 94% covered, and everywhere in Ireland will see greater than 90%. The further north you are, the better. (There are no cheap flights left to the Faroes – I checked )

Then, two weeks later on the other side of the moons orbit, the earth’s shadow eclipses the moon on April 4th. Unfortunately we are on the wrong side of the planet to see any of this. But, five days after the autumnal equinox, in the small hours of September 28th, we get to see a total lunar eclipse of our own.

For the March 20th solar eclipse, the times for Dublin are: eclipse begins 08:24, maximum eclipse 09:29, eclipse ends 10:37. At the time of maximum, the sun will reduce to a small crescent and the sky will noticeably darken, though not enough to see stars.

NEVER LOOK DIRECTLY AT THE SUN, even at the maximum of a partial eclipse, with the naked eye, binoculars, telescopes, sunglasses, or through beer bottles or other contraption you have dreamed up. Infrared light from the sun can damage your retina even if you feel you are not being dazzled. Your best bet is to use a small telescope or binoculars on a stand to project an image onto a white card (taking care not to burn yourself or start a fire), or make your own pinhole camera. A cheap natural version of the pinhole camera is to stand under leafy trees and see multiple images of the crescent sun on the ground, diffracted through the gaps in the leaves… there won’t be too many leaves around on March 20th though. WARNING: Looking directly at the sun through a telescope or binoculars will cause immediate eye damage.

(Ok, that’s the standard hysterical warning. At the max of the eclipse you can throw a quick glance sunward, just like you can accidentally glance at the full sun in the sky without going blind. But NO PROLONGED STARING and DEFINITELY NO BINOCS OR SCOPES: THAT STUFF ABOUT PERMANENT DAMAGE IS NOT A JOKE. I also can’t vouch for any eclipse glasses on Amazon that might arrive on time if you were to order now ).

#389

Thanks

I also found a partial answer here

eso.org/public/usa/news/eso1508/

#390

Thanks C&S, that article’s very interesting. I see the research team leader is a Paddy with a PhD from UCD who’s become a high flyer at the Niels Bohr Institute. Impressive CV too.

#391

Amateur Cosmology

There’s an impression that cosmology is the domain of irredeemably nerdy boffins who can’t eat dinner without scribbling the field equations of General Relativity on their napkins, and who speak a language the rest of us cannot comprehend. Well, it’s true that GR is a bit complicated, and some of the additional more speculative theories can be hard to follow. (I’m speaking for myself here … it may all be a piece of cake to you).

But that’s not to say that you can’t do a bit of amateur cosmology yourself. What I mean to say is that from the comfort of your own armchair, you can look back in time and observe galaxies from the early universe that nobody else has looked at, and draw significant conclusions about the reality of cosmic expansion. Here I’m going to walk you (at a fast trot ) through your very own observation of the Lyman alpha forest (de wha’? … don’t worry, all will be explained if you have the patience to bear with me), and see what it says about the evolution of the universe.

Now, you’re going to need to be able to interpret a spectrum, and for that you need to understand a little about spectroscopy, which in turn means you need a smattering of atomic physics … so it’s good that you were paying attention here and here. For an ultra-quick refresher, electrons orbit the nucleus of an atom like satellites orbiting the earth (not really, but let’s pretend), and can move between higher and lower orbits by taking discrete sized jumps corresponding to exact quantities of energy which we see as colours in a spectrum. Here’s a schematic of different possible orbits, labelled n=1,2,3… , for the single electron in a hydrogen atom (fyi, n stands for a fancy thing called the principal quantum number, but you can ignore that):

Each possible jump between orbit levels is associated with a particular amount of energy, corresponding to a particular wavelength of light (i.e. a particular colour), and you can see this in the diagram, e.g. the very first jump from n=1 to n=2 is labelled with “122 nm”. That means a wavelength of 122 nanometres, or billionths of a metre. If an electron (in a hydrogen atom) jumps down from n=2 to n=1 it emits a photon of light of that wavelength, and if it jumps up from n=1 to n=2 it must absorb a photon of that wavelength.

I need to also mention another unit, the Ångstrom, which is simply one tenth of a nanometre. So 122 nanometres is 1220 Ångstrom. As you see in the picture, the set of jumps from n=1 are collectively called the Lyman series, and since 122 nm is the first one, it is commonly referred to as Lyman alpha (or Lyα for short). Actually, its more precise wavelength is 1216 Ångstrom. Elsewhere you can see that “Balmer alpha” is 656 nm … but for historical reasons we just call this Hydrogen alpha (or Hα).

Ok, lets look at a star. Amazingly, the light we receive from these giant objects is strongly affected by the physics of tiny atoms we described above. The particular star we will look at is in a fairly empty and non-descript part of the sky between the constellations of Orion, Taurus, and Cetus. At the moment, it leads Orion in the sky, setting before it about 11pm. Now, you’re never going to see this star with the naked eye – it’s nearly a million times too faint. But the 2.5 metre telescope that took this picture not only photographed it, but passed it’s light through a prism-like device to split it into a spectrum. The following pictures show you where this star will be around 8.30pm tonight, followed by an actual photo (that fuzzy blue thing between the crosshairs), and then the star’s spectrum.

I’m going to explain the spectrum in some detail, because it’s not nearly as scary as it looks. Remember, we’ve split up the starlight into its constituent colours, and now we have graphed the amount of energy in the starlight at each different colour. Ok, you don’t see colours on this graph … instead you see wavelengths in Ångstroms marked along the horizontal axis. But these are just what our eyes see as colours – everything to the right of 6000 Å is red, everything to the left of 4000 Å is blue or violet, everything in between is green or yellow or orange. Remember the Hα line from our atomic model earlier? It’s at 656 nm or 6560 Å, so we expect it to be red. Actually, Hα is the type of light given off by hot hydrogen in star forming regions like the Orion nebula, and you may remember in an earlier article on star formation I showed a photo of this – in large part it does indeed glow a rose red colour. The rose-coloured chromosphere of the sun (also pictured in a previous article) that we see during solar eclipses is also glowing in hydrogen alpha light.

Let’s look at the spectrum itself. It’s the very wiggly, noisy black line (because of uncertainties due to this star’s faintness). Superimposed as a red line is a computer-interpolated cleaner version of the spectrum. It is telling us how much energy there is in the starlight at each wavelength marked on the horizontal axis. You can see it’s fairly flat along most of its length but with some significant pointed dips, and a general slope up toward the left. Lets get the explanation for that slope out of the way first. Bluer colours are toward the left, redder ones toward the right. Look at the photo of the star. It’s blue, indicating a very hot star. That means there is more power in its spectrum toward the blue end. That’s why the spectrum slopes upward toward that end.

Now look at the units on the vertical axis. They look scary but we can interpret them easily. A “flux” in this context is just an amount of power received over an area. An example that you will be well acquainted with if you’ve ever taken an interest in solar power is the solar flux at the surface of the earth. It’s about 1,350 Watts per square meter. That is, if you had a perfectly efficient solar panel with an area of one square meter, you’d get 1,350 Watts of power from it. Power is a measure of energy per unit time, and one Watt is the same as one Joule (of energy) per second. (You need about 500k Joules to boil a full kettle). So we could also rate our solar panel in Joules per second per metre squared.

Ok, for our flux from the star we are using almost exactly analogous units, but a slightly older system of ergs (of energy) per centimetre squared per second. This is the same as thousandths of Watts per metre squared. You also see that this measurement is per Ångstrom. Whereas for the sun’s energy on our solar panel we measure the total power across all the different wavelengths in the sunlight, in our star spectrum we measure the power falling on the telescope’s mirror in each Ångstrom wavelength interval. It’s as if we had a different solar panel for each possible colour of the rainbow, and we measured the sun’s power for each different one – each panel picking up only its relevant fraction of the power across the whole spectrum.

Also note that the units on our vertical axis are scaled by ten to the power of minus seventeen; the visible light from this star is about a million trillion times fainter than the sun. If you ever looked closely at a fly flexing his leg, he used more energy in that movement than the total energy from this star that would fall on our entire 2.5 metre telescope mirror if we were to point it at the star for a month! (or on the pupils of your eyes if you were to look at the star for a million years!)

Now let’s turn our attention to some other things that the computer has filled in for us. You can see vertical lines labelled with cryptic names, blue ones above and red ones below. They’re not all cryptic though … we should be able to recognise the Hα line marked in red, just where it should be, above the 6563 Å (656 nm) point on the horizontal axis. Look at where the vertical line intersects the computer-interpolated red spectrum line. There’s a dip in the spectrum right there. This is perfectly normal for a star. It has a hot surface with a cooler atmosphere, mostly made from hydrogen gas. The cooler hydrogen gas is absorbing some of the energy from the starlight and using it to pump electrons up from the n=2 to n=3 level, so this dip in the power of our spectrum is the Hα absorption line.

What else can we recognise? There are many more absorption lines in this spectrum, which again is normal for stars. Stars have absorption spectra because, as I said, we see their bright surface light through a cooler atmosphere. You might be able to guess that the Hβ line at 4861 Å (486 nm) is the second line in the Balmer series from the atomic model earlier, corresponding to a jump between the n=2 and n=4 electronic orbits. Once again we see a dip in our spectrum at this wavelength. If you know your Greek alphabet, you’ll be able to find other Balmer lines. I will just mention in passing that some other lines correspond to elements other than hydrogen … I can see sodium, calcium and magnesium for instance, and they give me telltale clues about what type of star this is, and how hot it is. Some of the lines labelled in blue can tell me how big it is. This is the sort of amazing detail that a spectrum can reveal about these far away objects, even though an ordinary photograph through the most powerful telescope only reveals a pinprick of light.

What you won’t see in this spectrum are things like Lyman alpha (Lyα) at 1216 Å (122 nm). Our spectrum simply doesn’t extend that far to the left. 122 nm is in the ultraviolet part of the spectrum that we can’t see, and our telescope and spectrograph aren’t sensitive to. Ultraviolet (UV) is more energetic (shorter wavelength) than visible light, and indeed the Lyα line marks the boundary between the very energetic far UV and extreme UV.

I hope you’ve stayed with me thus far, because in the next article we’re going to look at more spectra … but this time we are going to venture far out and long ago into the expanding universe, and we will turn our attention from stars to galaxies!

#392

Interactive gravity.

codepen.io/akm2/full/rHIsa

#393

cnet.com/news/8-possible-exp … net-ceres/

#394

The Dawn mission to Ceres is going to be every bit as exciting as the Rosetta mission to comet 67P … but without the plot twist of the little lost lander.

#395

Cosmeolaíocht Amaitéarach (Cuid a dó)

In Part I we looked at how electrons jumping between different orbits in an atom can emit or absorb photons of light of very specific colours/wavelengths. We took a particular interest in the hydrogen spectral series. Next we looked at how starlight (or any light) can be split up into its constituent colours to give a spectrum. The spectrum shows the flux (i.e. power received at detector surface) at each wavelength. The spectrum has an overall shape (called the continuum) which corresponds to the object’s colour, and may have emission or absorption features in the shape of spikes or notches. The continuum is produced by a body in thermal equilibrium (like a hot poker or a star’s surface), and the emission and absorption lines are most commonly produced by a rarefied gas (as in a vapor street lamp or a star’s atmosphere).

Now we take a look at a galaxy spectrum. First, though, I want to show you how you can investigate your own spectra. The Sloan Digital Sky Survey was twenty years in the planning, and so far has been fifteen years in the execution. I could wax lyrical about its amazing technology and equally amazing achievements, but there are umpteen science papers about it on the web if you want the technicalities, and for a very quick introduction one can do no better than this video:

That video is ten years old, and the equipment has been upgraded and various new surveys undertaken, but the basic techniques remain the same. The fantastic thing about the SDSS is that every scrap of data ever collected has been made freely available to scientists and the public alike. You can go to the sdss3.org/ website to start reading about all aspects of the survey.

Then go to the sky image browser. This lets you see the sky, exactly as photographed by the 2.5m SDSS telescope at Apache Point, New Mexico. You can pan around the sky, or type in celestial coordinates you want to go to. The default location when you open up the browser is a part of the sky near the head of the snake in the constellation of Hydra. I should point out that the SDSS 2.5m telescope is not a large telescope as modern scopes go, but its certainly not small either. It also doesn’t take different length exposures – it is designed to continuously “drift scan” the sky in one infinitely long continuous exposure. So the range of object brightnesses (or magnitudes) it can record is fixed. If you use a sky programme like Stellarium, it shows you stars down to magnitude 9 or 10. The SDSS detectors saturate at about 100 times fainter than that. In short, you are not going to see any familiar stars here – the very brightest is 10,000 times dimmer than you can see with the naked eye.

Ok, on the sky image browser, there are options to zoom and pan, and to search. However, without moving from the default location, click the “Objects with spectra” check box under “Drawing Options”. This places a red rectangle around each object for which the SDSS captured a spectrum. (Remember, that’s only the “interesting” objects for which they drilled a hole in one of those aluminium plates you saw in the video, so it could be fibre-optically wired to the SDSS spectrograph). Up near the top of the image you’ll see two red rectangles close together, one of them around a brightish blue star. Click on that. A thumbnail of its spectrum appears on the right. You also see coordinates (RA and Dec), a type (STAR, GALAXY, QSO=quasar]), and five numbers labelled u,g,r,i,z. You’ll come across these a lot – these are the magnitudes (brightness) of this object in the five colour bands that the SDSS cameras use: u = violet/ near-ultraviolet, g = green (centre of the visible spectrum), r = red, i = near-infrared (NIR), z = mid-infrared (MIR). The range of the (modern) SDSS spectrograph is about 3500 Å to 11500 Å, from near-ultraviolet through the visible spectrum to mid-infrared.

Ok, click on the explore link for this object (or click here). You can garner various information about the object, but we are interested in the spectrum. Clicking on the spectrum picture gives you a bigger version, but the Interactive Spectrum is more interesting as it lets you select horizontal and vertical sections of the spectrum to zoom in on. For this star, you once again see a “blue bump” around the 4000 Å region and you should be able to identify some of the prominent hydrogen absorption lines – if you scroll down you’ll see a list of their wavelengths. The Balmer series are the ones labelled H_alpha, H_beta, H_gamma, H_delta, H_epsilon.

Unfortunately I can’t go into all the different SDSS tools available to you, but feel free to figure them out for yourself. There are query pages for finding objects by type, including quasars (QSO type). You can even use SQL if you know that language, and have studied the SDSS database schema. And as well as viewing spectra interactively you can download the raw data for playing with in Excel or other graphing tools. There are many, many possibilities.

Back, now, to the galaxy spectrum we were promised. Here’s one I prepared earlier. Go ahead and click its interactive spectrum. It’s quite a faint spectrum but you should be able to identify weak absorption lines for Hα, Hβ and so on. These look like stellar absorption spectra and, indeed, that’s just what they are. A galaxy’s spectrum is just a sum of the spectra of its stars.

I should be careful to point out that you may not be looking at the whole galaxy. Galaxies are large extended objects, and even very far away ones are not points of light, like stars. To give you an idea of scale, suppose you scaled our sun down to a ball one inch across. (The earth would be about half the size of a baby pea). The nearest star to us on this scale would be 400 miles away! The distances between stars are vast compared to the sizes of stars. But suppose we scaled our galaxy down to a one inch disc (on which scale the sun would be about the size of a single atom!). The nearest major galaxy to us on this scale would be only two feet away! Galaxies are large compared to the distances between them. To be sure, there are also galaxies some miles away on this scale, and huge voids with no galaxies at all in them. But keep in mind when you look at a galaxy spectrum, you may be looking at older, redder stars in its central bulge, or perhaps younger bluer stars in its spiral arms. (Not all galaxies are spirals, of course).

Ok, take a look again at the Hα line in our galaxy spectrum. There’s something wrong with it. We expect it to be at 6563 Å, but here we see it at nearly 8000 Å. What’s gone wrong? Nothing – this is the famous cosmological red shift discovered by Edwin Hubble in 1929. The universe is expanding, and the space between the galaxies is expanding. As it expands, light that has been travelling to us from distant galaxies is stretched out to longer wavelengths. The portion of the red light from this particular galaxy that set out to us at 6563 Å (its emitted or rest wavelength) has been stretched to an infrared wavelength of 7914 Å. All the other lines have also been shifted toward the red end of the spectrum, by a factor of 1.206. Subtract one from the multiplicative factor and you get the value denoted as z, the red shift of the galaxy … 0.206 for this particular one. (Apologies that z is the same letter as was used for measuring a star’s metallicity in a previous article – they are completely different things: the astronomers just ran out of letters to use). Whenever you want to know the factor by which spectral lines will move, just look up the z value (it’s on the explore and spectrum pages we’ve been looking at). The factor to multiply by is 1 + z.

Now, a point that should be fairly obvious is that, since the SDSS detectors only operate within a particular range of wavelengths, and the spectrum of our target object has been shifted to the right compared to its emitted wavelength, we’re going to lose some of that emitted spectrum off the end. By the same token, we’re going to get some new parts of the spectrum appearing on the left of our graph. In effect, we are looking through a window of wavelengths limited by the SDSS capabilities, but as we look at further away objects, the shorter wavelength parts of their emitted spectrum slide into view in our window. This gives us a fascinating way to study parts of the electromagnetic spectrum that we otherwise could not see, or that we would have to launch UV and X-ray telescopes into space in order to see.

The next article will be about some of the things that this enables us to find out.

#396

Astronomers Watch a Supernova explode over and over again.

nytimes.com/2015/03/06/science/astronomers-observe-supernova-and-find-theyre-watching-reruns.html

Truly fascinating.
This is the ***same ***supernova.

#397

Cool!
I’ve got some techie stuff on Einstein crosses somewhere that I’ll dig out for a future article.