[Jim Lattis, Director, Space Place, University of Wisconsin-Madison]

Well, good evening and welcome to U.W. Space Place. This is one of our guest speaker nights, and I would like to introduce Professor Ed Churchwell of the U.W.-Madison Astronomy Department. Recently retired from that department, but he came to Madison from Germany from the Max Planck Institute in – in 1977 and has been here ever since, during which time he’s done a lot of really important research here, primarily as a – well, I shouldn’t say primarily – as a radio astronomer, certainly, but here at Space Place you can see a lot of exhibits with Ed’s fingerprints on them because he was very involved with the Spitzer Space Telescope and a project called G.L.I.M.P.S.E., milking – mapping the Milky Way, the plane of the Milky Way galaxy.

So, I’m really pleased to introduce Ed tonight. He’s going to give us a – a talk on, basically, cosmology and how astronomers can figure out how the universe has evolved since its origin. So, let’s welcome Ed. Thank you.

[applause]

[Edward Churchwell, Professor Emeritus, Department of Astronomy, University of Wisconsin-Madison]

Okay. Thanks. Thanks, Jim.

Well, thank you very much for coming out, particularly on a rainy night like tonight. This is a topic that I could spend a whole semester on, okay? So, I’m going to try to distill some of the main components of what we have at our disposal to understand how the entire universe, in fact, has changed from the – an instant after it was born until the present time and on into predict the future. So, hang on to your hats because you’re going to go through a fantastic journey here. And please, if you have questions, interrupt, and I’ll be glad to answer them if I can as we – as we go along.

Okay, Mark Twain once made this famous statement.

[slide featuring the follow quote from Mark Twain]

“There’s something fascinating about science. One gets such wholesale returns of conjecture out of such trifling investment of fact.”

[laughter]

[slide animates on the bullet point that this is – an apt description of cosmology in the first half of the 20th century]

And that was certainly an apt description of cosmology in the first half of the 20th century.

[slide animates on the following statement – The idea that we can determine the properties of the universe a tiny, tiny fraction of a second after its birth and predict its evolution into the future sounds presumptuous in the extreme, but this is exactly what cosmology attempts to do!]

The idea that we can determine the properties of – of the universe as a whole, a tiny, tiny fraction after its birth, is presumptuous to say – say the least. But, in fact, that’s precisely what cosmology tries to do.

[new slide with the heading – Cosmological Parameters and the statement – The history of the universe depends on the determination of a surprisingly small number of quantities of H0, Omega r 0, Omega m 0, Omega Lambda 0, q 0 and T0 and their dependence on time.]

So, surprisingly, the history of the universe can be described based on six quantities that – that can actually be determined from observations, luckily. And I’m just going to kind of point out quickly what these are. H-naught is what’s called Hubble’s constant. It basically is a measure of how fast galaxies or clusters of galaxies –

[Edward Churchwell, on-camera]

– are receding away from us as a function of their distance from us. And that’s a fundamental quantity, and it’s one that’s easily observed. In fact, it was first observed in 1929.

The omegas here, one is omega sub-R, it means for radiation. It’s the fraction of energy density in the universe in the form of radiation or photons. And the naught – the naught here means that these are quantities that we measure at the present time, okay.

[return to the Cosmological Parameters slide]

So, the naught means they’re – theyre – theyre quantities we measure today, but if we can measure them today, we can project them back at any other time. And that’s a key issue, mainly because if we can establish that the universe was in thermal equilibrium, then we can determine –

[Edward Churchwell, on-cam]

– all these – all these parameters at any given time.

The omega M here is the fraction of energy density in the universe in the form of matter, okay. The omega – the omega sub lambda is the fraction of energy density in what we call dark energy. And I’ll come back to all of these in a moment.

[return to the Cosmological Parameters slide]

Q-naught is a – is called the deceleration parameters. So, it’s – its a parameter that determines how – how – what the rate of slowing down of the expansion of the universe is.

[Edward Churchwell, on-cam]

And basically, it’s not an independent parameter. It really depends on omega-naught. So, if we can determine omega-naught, which is the sum of all of these – of all of these three omegas I was talking about here, then we also know Q-naught.

[return to the Cosmological Parameters slide]

And T-naught is the temperature of the universe as measured today. So, lets –

[the Cosmological Parameters slide animates on the following statement – These quantities can be measured and their dependence on time calculated via well-established principles]

– let me move on now.

[new slide titled, Observational Underpinnings, with seven bullet points – 1) Universal expansion (Hubbles Law, 1929)= H-naught, 2) Abundances of light elements – primarily D and He – 1950s = hot, dense past, 3) Dark Matter – 1933,1939,1970s – more in the universe than meets the eye, 4) The Cosmic Microwave Background (CMB) – 1965, 1990 = T-naught thermal equilibrium, 5) Large Scale distribution of Matter – 1990s = nature of gravity and expansion, 6) Recession velocities of Type 1a Supernovae – 1998 & 1999 (2 groups – Reiss et. al. and Perlmutter et. al. = acceleration of the universe, and 7) Galaxy clusters, gravitational lensing, and early propagation]

So, the observational underpinnings. There are about six kinds of observations that have been made that actually determines these properties and tells us how to go about using those properties to say what the universe was like at any given time.

One is the – what’s called the universal expansion, or Hubble’s Law, which was first published in 1929. He got H-naught a little wrong, but he – he certainly understood the – the principle. The other is the abundance of the light elements of primarily deuterium and helium, and that was something that first kind of came to light in the 1950s.

[beeping]

And the others, dark matter –

[beeping]

– the detection of dark matter – is that me? –

[Jim Lattis, off-camera]

No, youre good. Just the stand is [unintelligible]

[Edward Churchwell, off-camera]

Oh, sorry. [laughs] The light’s in my eyes. [laughs] I’m trying to move out of – out of the light.

So, dark matter was first actually detected in 1933 by a Swiss astronomer who was working at Caltech at the time. And he was – was looking at the motions –

[Edward Churchwell, on-camera]

– of galaxies and a cluster of galaxies and realized that they were moving a lot faster than they should be moving based on the amount of light that they were producing. In 1939, another astronomer I’ll get to in a moment, also found roughly a – a similar sort of thing, but – but basically by measuring the rotation rates of spiral galaxies. And he found they were more – they were rotating much faster than they should. And then, in the 1970s –

[return to the Observational Underpinnings slide]

– Vera Rubin did a very detailed study of the rotation of galaxies and basically found that the only way you could understand the rotation rates is if there was more – more matter –

[Edward Churchwell, on-camera]

– there than you could actually see from the light that’s produced by stars. And Im – I’m gonna – Im gonna kind of to go through these a little bit more in a moment. So, the large scale distribution of matter, oh, the micro – the cos – the cosmic microwave background – let’s come back to that – that’s an important one –

[return to the Observational Underpinnings slide]

– was first detected almost by accident in 1965 by two radio – radio engineers, actually, at Bell Labs. And basically, what they measured was the temperature of the universe today – at – at present time. Then the large-scale distribution of matter, that is how clumpy is the universe – the matter distribution on the universe?

[Edward Churchwell, on-camera]

And it turns out that it’s not smooth, at least on a – on a – what we call from an astronomical point of view, on a small scale. Large enough scale and it looks pretty smooth, but it ain’t necessarily so.

The recession velocities –

[return to the Observational Underpinnings slide]

– of type 1a supernovae is used – those are standard candles. There – theres objects that have well determined intrinsic brightness, and so if you have that –

[Edward Churchwell, on-camera]

– you have a standard candle so that you can determine from the apparent brightness, and you can determine what the distance is.

Okay, so I’m gonna – those are the six, kind of, main observational underpinnings.

[slide titled, Universal Expansion, with the statement that observations show that the velocities of galaxies are speeding away from us with increasing speeds at greater distances (Hubble & Humason, 1929) and under this statement are three bullet points, v=H0d, were v is the velocity of recession with distance; H0 is the present value of Hubbles constant, and d is the distance; this is truly independent of the direction = isotropic; H0=v/d = the rate at which galaxies increase in recession velocity with increasing distance – the rate of expansion of space-time with distance – at the present time]

And I’m just going to quickly run through this. Universal expansion basically was done by Hubble and Humason in 1929, and they found that the velocity, the recession velocity that you can measure through spectroscopy is basically just proportional to the distance of – of the galaxies. And so, this H-naught is the key parameter here. So, if you just take the velocity which you can measure and you have to be able to measure the distance independently in some way, but there are ways to do that. And so, H-naught is this V over D. And so that’s – thats what was used to – to determine this.

[slide animates on a further bullet point that this result implies that the universe was much smaller in the past – i.e., mass densities were much greater]

This result, however, implies that the universe is much, much smaller in the past, i.e., the mass densities are much greater than they are today. Also, if the universe is expanding and you’re not changing the total energy in the galaxy –

[Edward Churchwell, on-camera]

– what happens when you take a – a gas under pressure, and you let the pressure off? What happens to the temperature of the gas?

[first audience member]

It drops.

[second audience member]

It goes down.

[Edward Churchwell]

It goes down, right? That’s exactly the same thing that happens in the universe. The universe started out as something incredibly dense and hot and massive and – and then it has been expanding ever since. At one point it was explosively expanding. And so, as it expands, it also is cooling down. And so that’s – thats a key – the key component of our understanding of the evolution of the universe.

Expansion here, you got to be careful. It doesn’t mean – it’s – its not like the motion of galaxies through space. It’s space itself that’s expanding, okay. So, it’s like the – the – the – there are several examples that have been given that people use, but one is if you have a loaf of bread with raisins in it, you start to cook the bread and the bread rises, the relative orientation – positions of the raisins don’t change. They – they – they only change in their separation, right? But – but the pattern doesn’t change. And this – this is the same sense we’re talking about for the universe, that – that it’s space-time is- that is expanding not the motion of galaxies through space. Think about that.

Okay, abundances to light elements.

[slide titled, Abundances of Light Elements (D, He, Li, & Be), and the bullet point – abundances of these elements can only be understood if the early universe was hot and dense enough to have been a thermonuclear furnace. That is, they could not have been formed in stars at there current abundances. This is also consistent with, and a confirmation of, the implications of Hubbles Law.]

Those elements, deuterium, helium, can only be understood if the early universe was hot and dense enough to have been a thermonuclear furnace, and therefore have cooked the hydrogen atoms and produced helium and deuterium. The stars can’t produce enough of this to – to – the – – the – the –

[Edward Churchwell, on-camera]

– of the amount that we actually see. Okay, so – and this is consistent with the idea –

[return to the Abundances of Light Elements slide]

– of the expansion of the universe. If you go back in time, the universe had to be hot and dense, and at some point, it had to go through a period where it was dense enough and hot enough to actually behave like a thermonuclear furnace. Okay.

[slide animates on a new bullet point – Alpher, Herman, & Gamov based on the realization that the very early universe was hot and dense enough to sustain nuclear reactions that would convert H to He, and possibly heavier elements in the 1950s. Alpher, Bethe, and Gamov also predicted that the current temperature of the universe should be approximately 5 Kelvin – a prescient insight at that time]

Alpher, Herman and Gamov, based on the realization that – of the universe was hot and dense enough to sustain nuclear reactions to convert hydrogen to helium and possibly heavier elements in the 1950s. And they actually even predicted roughly –

[Edward Churchwell, on-camera]

– what the – what the temperature of the universe should be today. And they – they estimated somewhere at five to seven degrees Kelvin. And you’ll see in a moment that – that we have a very precise number for that now, and it’s about three degrees.

Dark matter, Fritz Zwicky –

[slide titled, Dark Matter, with the bullet point that Fritz Zwicky in 1933 found that the velocities of galaxies in clusters of galaxies were too large for the mass indicated by their luminosities. He postulated that this was due to the presence of dark matter not detectable by his telescope]

– in 1933 found that the velocities of galaxies – I’ve already gone through this, so I’m going to zip through it fairly quickly. But, so that you see, though –

[slide animates on a new bullet point – Horace Babcock in 1939 found that spiral galaxies rotate faster than they should based on the amount of mass indicated by the light emitted by their constituent stars. He didnt comment on the cause of this]

– there were several players in this – in this area.

[slide animates on the final bullet point that Vera Rubin in the 1960s & 1970s did a systematic study of spiral galaxy rotation motions with distance from their centers (nuclei) and found that they systematically rotate much faster than they should based on the amount of mass indicated by the stars contained in them. Her work established the presence of dark matter in galaxies that does not emit or absorb electromagnetic energy (photons)]

Vera Rubin, in the ’60s and ’70s, basically did a – this systematic study of the rotation of spiral galaxies and, you know, concluded that – there – there’s a lot of unseen matter in the – in the universe that we don’t see –

[new slide titled, Schematic of Galaxy Rotation, featuring a graph with velocity on the y-axis and distance on the x-axis and with two lines, one – marked A is the predicted rotation curve of a spiral galaxy which increases dramatically at a short distance but then tappers off like half of a bell curve the longer the distance, and a second, marked B, which is the actual observations which indicates a spike at shorter distances and a straight line with no tapering at longer distances noting that dark matter can explain why the actual rotation curve doesnt tapper off as predicted]

– you know, with our telescopes. This is the idea, however, right here that if you – if you look at the rotation rate of spiral galaxies from the center of the galaxy outward in radius, what you would expect if – if all the matter was confined to the nucleus of the galaxy, it should go like the curve A that I show here. This is how the – the planets move around the sun. They go as one over R – one over R to the one-half power. But that’s not what happens. They actually – they – the rotation rate stays up here flat, much higher rates of rotation as you go outward than – than you would expect. And the only way that can happen is if you – you’re enclosing more and more matter as you – as you go further out in – in radius on these spiral galaxies.

[new slide titled, The Cosmic Microwave Background (CMB)]

Okay –

[slide animates on the bullet point – the radiation temperature of the universe today is 2.735 plus or minus .005 Kelvin. It is the same in all directions, and is essentially a perfect Planck Function (Blackbody = thermal equilibrium – a Maxwelling velocity distribution] according to Penzias and Wilson, 1965 and Cobe satellite around 1990)]

– cosmic microwave background. The radiation temperature of the universe today is about 2.735 plus or minus 0.005 Kelvin. We know that very precisely now –

[Edward Churchwell, on-camera]

– because basically the Planck, the C.O.B.E. – the C.O.B.E. satellite in the 1990s. But anyway, the fact is that you can look in this direction, this direction, any direction in the sky, and you get exactly the same temperature, or you get exactly what’s called the – the Planck function. It’s a function that tells you the brightness as a function of wavelength. And we know that very precisely, and I’m going to show you here.

[return to the Cosmic Microwave Background slide now with an added bullet point – The CMB is a picture of the universe when it was last opaque. Its observed temperature today is only 2.7 Kelvin due to its great red shift (at approximately 1100, when the universe was only approximately 280,000 years old). This is consistent with the implications of the expansion of the universe and the light element abundances. That is, all the observational evidence consistently indicate that the early universe was very, very dense and very hot and in complete thermal equilibrium, which allows us to track its properties (contents, temperature, density, etc.) with time]

What you – what – so – so, the cosmic microwave background is really a picture of the universe when it was about 280,000 years old, okay? This is – you can look back into time – into space, and at some point, you run into a wall where the universe becomes opaque.

[Edward Churchwell, on-camera]

You can’t see any deeper because it becomes opaque. And, that – at that point, if you – if you take a map of the sky, what you see is, you know, this microwave background that’s only about 2.73 degrees. But – at your – when its – when it was 280,000 years old, it was very, very hot. But today we see it for only three degrees because the – the stretching of space-time has stretched out the wavelength now to where it’s equivalent to about three degrees. Okay, so – And it also – oh, an important thing. That tells us, however, that the universe was in complete thermal equilibrium at that time. So, we can use this idea, you know, using basically basic physics, to trace back in time, or go forward if we wish.

[slide titled, The Cosmic Microwave Background – C.O.B.E. Results, featuring a graph with brightness on the y-axis and frequency/wavelength on the x-axis showing a modified bell-shaped curve with brightness maxing out at a frequency of 150 and then tapering off as frequency increases and showing that the predicted and observed graphic are the same line]

Here, just to let you – get an idea, though, of how precise a measurement this is. This is one of the most precise measurements you’re going to see in the – in the – in the natural world. The dots that you see in that curve are measurements from the C.O.B.E. satellite. The – the curve is a Planck function. It’s the theoretical function for the brightness of an object at a given temperature. Take a bar of iron and heat it up and it goes from very dull red to a bright red to a yellow to white. And that’s a Planck function, okay? Its – and it’s determined entirely by one quantity: the temperature. It doesn’t matter what it’s made out of, anything else, except the temperature. So that’s the – thats the curve of – of the brightness of the universe as seen today. And it peaks at about two microns. Or – or two millimeters, I’m sorry.

[new slide titled, CMB – The Universe at approximately 380,000 years old, featuring an oval image with primarily a green pixelated color that is broken up occasionally by an orangish-red pixelated color and a blue pixelated color]

Okay, so here’s a picture of the universe or of the cosmic microwave background when it was about 380,000 years old. And you see all these – these – the orange or reddish dots are a little bit warmer. The green, intermediate, and blue is a little cooler. But the amount of deviation from a – from a single temperature is tiny. It’s in the 10 to the minus four. One-ten-thousandth from top to bottom here. So, these are – are little ruf-ruffles which ultimately later on become – are going to become galaxies and things like that.

[new slide titled, View of the Universe with Age, featuring an illustration with three oval shaped diagrams on it flowing from top left to bottom right, where prior to the first diagram is the Big Bang, the first diagram is an orangish, swirling color labelled tiny fraction of a second, it flows into the second diagram is the previous illustration of the CMB in the prior slide labelled 380,000 years, that flows into the final diagram which is an illustration of the modern universe with galaxies in it and labelled 13.7 billion years]

So, view of the universe. I – I’ve just shown you one at about 380,000 years. Going back even further in the past, the universe is – is basically just a hot thermal cauldron of – of subnuclear particles. And then, if we go to present day, the universe looks like the bottom picture there. So, this is kind of a very quick thumbnail picture –

[Edward Churchwell, on-camera]

– of the – of the change in the universe. But a lot of changes have occurred, and I want to go through a few of those.

As I said –

[slide titled, The Distribution of Matter in the Local Universe, with two bullet points, the first – The residual matter (baryionic matter) and dark matter develop into a web of intersecting filaments where galaxies form mainly at the intersections of the filaments. The filaments are a natural consequence of gravity in an expanding and cooling universe. Thus the observed filamentary structure of the matter in the Universe and permits us to understand the origin of galaxies, the second – the cosmological principle applies over a substantially larger volume than previously thought]

– one of the other things that we want to know is the distribution of matter. How – how clumpy is the matter that’s distributed in the universe, and on what scale? And, basically, ultimately what happens here, as – as we’ll see, is if – if you’re thinking about the universe as space-time expanding with time and gravity operating with this expansion, what happens? Well, a lot of simulations have been done, and – and basically what should happen is that matter starts to form into a filamentary structures and things like that. And so, its not – it will not be smooth. Itll – matter tends to want to gather together with gravity. And so, we want to know if that in fact is what happens.

[new slide titled, Observed Large Scale Structure of Matter, featuring and illustration of three blue oriental fan-shaped sets of data displayed horizontally from three different surveys with galaxies marked by yellow dots. The y-axis indicates light-years and the x-axis indicates the number of galaxies]

And these are observations now of galaxies. And there were three surveys here. This inner one had a fairly wide – wide angle. And all of these little yellow dots here are galaxies. And you see that they’re not smoothly distributed. They – they are kind of filamentary. And there’s a second survey that was a little bit narrower piece of sky but going deeper, and you see that, again, matter is not uniformly distributed. Galaxies are not uniformly distributed. They’re in kind of filamentary structures.

[Edward Churchwell, on-camera]

And then an even deeper survey where you – you start to run out of galaxies out there, not because there aren’t any there but because it’s a sensitivity issue that you can’t detect these very distant galaxies so easily. But you see again –

[return to the Observed Large Scale Structure of Matter slide with the galaxies graphed]

-no matter where you look, the distribution of galaxies is not smooth. It’s – its filamentary. But, of course, if you add this up over a long distance, a large enough volume, the universe does look kind of uniform. And that’s important because all the relationships that we use make two assumptions. One, that the universe is isotropic. That means it doesn’t matter which direction you look; it looks the same.

[Edward Churchwell, on-camera]

And it’s uniform. It is uniform density. And, as I say, you’ve got to now think, however, cosmically in order for that to be true.

This is a simulation –

[slide titled, Simulated Large Scale Structure, noting that gravity gathers matter into massive filaments and clumps – providing places of star, galaxy, and galaxy cluster formation, and featuring an illustration what is perceived to be an orange/red/black primordial plasma-like structures]

– using basically just the idea of expansion of the universe and cooling and – and gravity, and this is what simulations tell you – you – you should see.

[new slide titled, Tools Needed to Understand the Evolution of the Universe]

Okay, what are the tools needed to understand the evolution of the universe?

[slide animates on first bullet point – application of well-established physics (nuclear, high energy, relativity, thermodynamics, etc.)]

Well, applications of well-established physics, okay? Nuclear physics, high energy physics, relativity, thermal dynamics, and so forth. But none of these are – are things that are – that are mysterious or not understood, okay. So, we can understand the universe using pretty straightforward physics.

[slide animates on a new bullet point under the application bullet point – physics of annihilation and pair production]

One of the key ones here is this physics – is what’s called physics of annihilation and pair production.

[slide animates on a new bullet point under the physics of annihilation bullet point – that annihilation is independent of temperature]

An annihilation, it turns out, is independent of temperature. It doesn’t matter. Once the temperature becomes low enough, the – some particles will annihilate with their antimatter particles. Matter and antimatter particles will – will annihilate each other. And if there were equal numbers, we would be – not have anything left in the universe. It would all be in the form of photons.

[new bullet point animates on underneath the annihilation bullet point – pair production is strongly dependent on temperature; T is equal to or greater than m0c2/k approximately 6 times ten to the thirty-sixth m0 (gm)]

Pair production is strongly dependent on temperature. Pair production is when you have a – a – two photons –

[Edward Churchwell, on-camera]

– that collide with each other, and if they’re energetic enough, they will produce a – a matter and an antimatter pair particle, okay. But they’ve got to be hot enough – the – the particles have to be moving fast enough or have to be hot enough to be able to – to have the equivalent energy of the mass particles that’s produced. And so, basically, the temperature has to be greater than –

[return to the Tools Needed to Understand the Evolution of the Universe slide]

– it’s proportional – greater than the sum factor times the mass of – of the particle that’s actually produced. That’s the – the key. So, one of these is incredibly dependent on temperature. The other is independent of temperature. And so, this makes for an interesting evolution of the universe.

[slide animates on a new bullet point underneath the pair production bullet point – threshold temperatures – T(neutrons) is greater than or equal to 1.090 time ten to the thirteenth Kelvins; T(protons) is greater than or equal to 1.088 time 10 to the thirteenth Kelvin; T(electrons) is greater than or equal to 5.93 times ten to the nineth Kelvin]

So, these threshold temperatures, the neutrons can be produced by colliding two photons together that as long as the temperature is greater than about 10 to the 13th Kelvin. Protons are – are a little less massive, and they have to be hotter than about 1.088 times 10 to the 13. So, there’s not much difference, but there is difference in mass and therefore the difference in the threshold temperatures. Electrons, of course, are much less massive. And, of course, their – their threshold temperature is substantially lower. It’s about six times 10 to the 9th or six billion Kelvin.

[slide animates on a new bullet point beneath the Physics of Annihilation bullet point – Cosmic temperature today implies that temperature in the past when the universe was much smaller and denser was also much hotter. Temperature can be inferred back to a very, very small fraction of a second (Talpha1 divided by t)]

Okay. So, the cosmic temperature today implies the temperature in the past when the universe was much smaller –

[Edward Churchwell, on-camera]

– and denser, much hotter and much denser than – than today of course. And you can infer the temperature back to a very, very small fraction of a second. So, temperature is basically just proportional to – it should be one over the square root of the temperature here, of time, of age of the universe. So, it should be a square root over the t, the little t.

Okay, so the equation of motion that includes all of the constituents of the universe. That means baryons and –

[return to the Tools Needed to Understand the Evolution of the Universe slide, now with a new bullet point underneath the cosmic temperature bullet point – Solution of the equation of motion that includes all the constituents of the universe – matter (baryons and dark matter); radiation; neutrinos, dark energy, and geometry of S-T]

– dark matter, radiation, neutrinos, dark energy and so forth, can be solved, and we’ll come to that in a minute.

[new slide titled, Evolution of the Universe – ten to the minus forty-third of a second]

So, let’s talk about the evolution of the universe. In the earliest point at which we can say anything at all about the universe is about 10 to the minus 43rd of a second. Okay, so that’s like –

[Edward Churchwell, on-camera]

– a dot with 43 zeros after it and one, okay. So, it’s a tiny, tiny fraction of the universe. But believe it or not, we think we know roughly what the universe was like at that point. We can’t say much beyond that. The reason is because we’re not sure what – how the forces of nature, for example, were. I mean, the – the presumption is that the four forces of nature were basically just a single force at this time. We don’t know how those subnuclear particles would react with each other under these temperature and density and with the knowledge of – for example, we don’t know enough about gravity at that – at that stage. So, that’s the earliest moment we can determine –

[return to the Evolution of the Universe slide now with two bullet points – 1) The earliest moment we can determine the properties of the universe – the Planck epoch or hadron epoch, and 2) Temperature and density are so high that we are not sure how particles interact with each other at temperatures and densities higher than this time. Even the four forces (electromagnetism, gravity, strong, weak) are believed to be a single, unified force, although by this time gravity may have already separated out as a separate force. The universe is populated on by fundamental particles and their antimatter particles. Electrons, neutrinos, quarks, their antiparticles, and photons and gluons]

– the properties, and we’re not sure how particles interact with each other.

Okay, but we can say one thing: that the universe at this instant had to be made entirely of fundamental particles. That is, the particles that can’t be reduced any more. And – and what are those? Well, they’re electrons, neutrinos, quarks, and they’re antiparticles and photons and possibly gluons.

[new slide under the Time = Ten to the minus thirty two seconds heading with the subheading – The Period of Inflation]

Okay, we’re going to now go to a time of about 10 to the minus 32 seconds. This is what’s called a period of inflation.

[slide animates on the bullet point that the temperature is approximately 10 to the twenty first degrees Kelvin]

Temperature is about 10 to the 21 Kelvin.

[slide animates a new bullet point underneath the Temperature bullet point – The temperature is too high for protons and neutrons to be stable. If formed, collisions with other particles are so violent that they are immediately broken apart into their constituent quarks]

The temperature here is too high for protons and neutrons to be stable. So, what you have – these free quarks that are – that are moving around in – in space at very high velocities and very high densities. And if you form a neutron or a proton, they’d be broken down immediately because they’re being banged into each other at such high energies.

[the slide animates on a new bullet point – The universe has cooled more than a factor of a million and increased in size by tens of orders of magnitude in the period 10 to the minus thirty five to 10 to the minus 32 seconds (period of inflation)]

The universe has cooled from the time that I told you, 10 to the 43 – minus 43 seconds, by about a factor of a million or so and increased by size by tens of orders of magnitude. Now, the – the universe at this time has basically explosively – it’s started explosively expand. Nobody understands exactly why this is true – true, but it’s called the period of inflation. And the universe expands at a rate that is faster than the speed of light. Now, you may ask, how can that be? Well, remember, we’re not talking about motion of particles through space. We’re talking about space itself. The rate at which space can expand or contract is not controlled by – by the speed of light. Okay, by the theory of relativity. And so, the universe must have gone through this period of what’s called inflation where it may have expanded by, you know, like, 25 orders of magnitude in, you know, a tiny, tiny fraction of a second.

[slide animates on a new bullet point – Quarks and anti-quarks are beginning to be bound into protons and neutrons and their anti-particles, and as a consequence do not undergo freeze-out as individual quarks. The precise time when this occurs is not well established, but certainly earlier than when protons and neutrons freeze out of the universe]

So, quarks and antiquarks are beginning – by the way, quarks are these particles – they’re very strange particles, but they are – they are fundamental particles as far as we can determine. And protons –

[Edward Churchwell, on-camera]

– and neutrons are made of combinations of three different quarks. There’s another set of – of particles called mesons that can be formed out of combinations of two quarks. And – but at this point, they have to be – they’re free – free particles. And that’s something that we’ll probably never see again in nature. Okay, so they’re beginning to be bound into protons and neutrons and their anti-particles. As a consequence, they now undergo what’s called freeze out. That is, they – they no longer are – are going through – they don’t go through a period of – of annihilation and pair production. The precise time when this occurs is not all that well established yet. And I’m not going to say anything more about quarks.

[slide titled, Time = 10 to the minus 12 seconds]

Now we’re going to jump down to 10 to the minus 12 of a second.

[slide animates on the bullet point – the Temperature is approximately 10 to the fifteenth degrees Kelvin ( a factor of a million cooler)]

Temperature now is about 10 to the 15th Kelvin. It’s a factor of a million cooler than –

[slide animates on a new bullet point under the Temperature bullet point – Quarks have combined to form protons, neutrons, and mesons, and their antiparticles, leaving no free quarks]

– that other time point. Quarks have combined to form protons and neutrons and mesons and their antiparticles, leaving no free quarks.

[slide animates on a new bullet point – Electrons, neutrons, protons, neutrinos, and their anti-particles are stable and freely interact with radiation (pair production and annihilation rates are equal)]

Electrons, neutrons, protons, neutrinos, and their anti-particles are stable and freely interact with radiation, and so you have pair production and annihilation that’s going on at – at exactly the same rates.

[Edward Churchwell, on-camera]

So, they’re annihilating each other and being formed through pair production at – at equal rates. And so, it’s a stable, completely – what’s the word I want to use here?

[Male audience member, off-camera]

Zero sum?

[Edward Churchwell, on-camera]

Well, zero sum would be a good – a good term.

[Second male audience member, off-camera]

Equilibrium?

[Edward Churchwell, on-camera]

They were in equilibrium. That’s what I really want to say.

Okay, so now we’re going to jump down to a –

[slide titled, Time = one hundredth of a second]

– hundredth of a second.

[slide animates on the bullet point that the temperature is approximately equal to 10 to the eleventh degrees Kelvin]

Temperature now is down to about 10 to the 11th Kelvin.

[slide animates on a new bullet point – This is below the threshold temperatures of neutron and proton pair production, and both are rapidly disappearing (freezing out from the universe. N(n) is approximately N(p)]

This is below the threshold temperatures of neutron and proton pair production. And so, they’re rapidly disappearing from the universe, okay? Because theyre – they no longer can be formed at – at – by pair production because the temperature is getting too low, but they can – they can annihilate each other. And they keep right on annihilating until there aren’t anymore. And so, protons and neutrons are disappearing. And the only thing that saves them is the fact that for some reason that we still don’t understand, there was just a tiny little bit more matter –

[Edward Churchwell, on-camera]

– particles than antimatter particles. And so, all the antimatter particles disappear, and you get a tiny residual of matter – neutrons and protons. It’s about one out of every billion particles, okay? So, this – this is a key issue. Now, the – so, these – what we call – when this happens is what we call freeze out. The universe – the universe has gotten too cold to make more of these particles out of – out of photons. And so, they freeze out from the universe. And to begin with, the number of neutrons is about equal to the number of protons, but neutrons, when they’re not in an atomic nucleus are unstable and will decay at something like 11.7 minutes or something like that. So, and – and it produces a proton and an electron when it decays. And so, you get – you start to see then a change in the neutron to proton ratio. And this is important because it – that will ultimately determine how much helium the universe will be able to make.

Okay, so the density of the universe is about 3.8 times 10 to the 9th -b

[return to the Time = one hundredth of a second slide with a new bullet point – The density of the universe is about 3.8 times 10 to the nineth grams per cubic centimeter (energy plus matter). If Mount Everest were made of matter this dense, it would destroy the Earth by tidal forces]

– grams per cubic centimeter at this time. And if Mt. Everest were made of matter this dense, it would actually just destroy the Earth by tidal forces.

[slide animates on two new bullet points – the first, the doubling of time for the universe is about .02 seconds, and the second, at this time, there are about 10 to the nineth protons, electrons, and neutrinos for every proton and every neutron]

Okay, the doubling time for the universe is about two-hundredths of a second. And at this time, there are about a billion photons, electrons, and neutrinos for every proton and every neutron. Okay?

[slide animates on a new bullet point – Size of the universe? Dont really know, but using the present Cosmic Microwave Background temperature (3 Kelvin) and the temperature at this time, we can say that it would be 3 times ten to the eleventh smaller than today]

The size of the universe we don’t really know, but we can say that the universe at this time was something like 300 billion times smaller at this time than it is today.

[new slide titled, Time = .11 seconds]

Okay –

[slide animates on the bullet point that the temperature is now approximately 3 times ten to the tenth degrees Kelvin]

– you go to a tenth of second, temperature now is reduced to about three times 10 to the 10th Kelvin.

[slide animates on a new bullet point – Contents of the universe = electrons, positrons, neutrinos, anti-neutrinos, and photons. Protons, neutrons and their antiparticles have essentially disappeared from the universe leaving behind only a tiny fraction (10 to the minus nineth) of their original number. Why were there more particles than antiparticles? A mystery!]

Contents of the universe now is electrons, positrons, and neutrinos, antineutrinos, and photons. Protons and neutrons and their antiparticles have essentially disappeared, leaving behind this tiny residual of about one billionth. So, for – you know, theyre less than – less than – one billionth of their – of their original number.

[slide animates on a new bullet point – All particles are in thermal equilibrium (have the same temperature) and temperature is above the threshold temperature of all particles except the small number of neutrons and protons]

Let me just move on. So, all particles are still in thermal equilibrium. They have the same temperature and temperatures as the above threshold temperature.

[slide animates on two new bullet points, the first – Energy density is proportional to T to the fourth power – density equivalent to approximately 3 times 10 to the seventh power grams per cubic centimeter, about 300 million times that of water, and the second – the doubling time of the universe is about two tenths of a second]

I’m just going to go through this fairly quickly. The doubling time of the universe now is about two-tenths of a second. So, you see that this –

[Edward Churchwell, on-camera]

– this – the universe is already slowing down in its doubling time or its rate of expansion. And it continues to slow down because the universe is expanding against its own gravity. And the temperature and matter density continue to decrease as – as – as time goes along. And the neutrons are starting to get out of –

[slide titled, Time = 1.09 seconds with the bullet point that the temperature is now approximately 10 to the tenth degrees Kelvin]

– out of equilibrium or out of – they’re getting fewer – fewer neutrons than protons with time.

Now, at about 10 to the 10th Kelvin at about one second –

[a new bullet point animates on the slide – This is about twice the threshold temperature of electrons and positrons and they are beginning to annihilate more rapidly than they can be created out of radiation]

– the threshold temperature of electrons and positrons now are beginning to annihilate more rapidly than they can be formed through pair production. And so, electrons are starting to disappear.

[slide animates on a new bullet point – It is still too hot for neutrons and protons to combine into atomic nuclei. Both protons and neutrons are only about 10 to the nineth of their original values (the fraction of matter particles over anti-matter particles]

It’s too – still too hot for neutrons and protons to combine into atomic nuclei. So, these are all just sub – subnuclear particles. Nuclei can’t hold together at these high temperatures and densities.

[slide animates on a new bullet point – Density is approximately 3.8 times 10 to the fifth grams per centimeter squared. At this density neutrinos are beginning to cease interacting with matter and radiation (cross sections too small)]

The density is about four times 10 to the five grams per cubic centimeter. Water has a density of about one gram per cubic centimeter, alright? So – so, this is three times 10 to the five times more – more dense than – than water.

[slide animates on a new bullet point – Doubling time is approximately two seconds]

Doubling time is about two seconds.

[slide animates on a new bullet point – The decreasing temperature has allowed the proton-neutron balance to become approximately 24% neutrons and approximately 76% protons]

A decrease in temperature has allowed the proton to neutron balance to become –

[Edward Churchwell, on-camera]

– 24% neutrons and 76% photons, or protons, I mean. Now we go to about –

[slide titled, Time = 13.8 seconds, with a bullet point noting the temperature is now approximately 3 times ten to the nineth degrees Kelvin]

– 13.8 seconds. The temperature is about three times 10 to the 9th K.

[new bullet point animates on the slide – This is below the threshold temperature for electron and positron production, and they are rapidly disappearing (via annihilation) from the universe. This adds 0.511 M.E.V. of energy/annihilation and slows the rate of cooling of the universe until all the positrons are gone, leaving a small residual of electrons (approximately ten to the ninth)]

Here we’re below the threshold temperature for electron and positron production. And so, they are rapidly freezing out as well. So, now the electrons are disappearing as well. And this does add some energy. Each annihilation adds about a half a M.E.V. of energy, which actually raises the temperature of the universe by a little bit.

[new bullet point animates on the slide – The neutrinos and antineutrinos have decoupled from the rest of the particles and do not share the heating by electron-positron annihilation and are therefore approximately eight percent cooler than the rest of the universe]

Okay. Neutrinos and antineutrinos at this time are more or less decoupled from the rest of the universe. They no longer interact with – with normal matter. And, you know, there’s about a billion of these neutrinos that goes through my fingernail while I’m speaking, you know, at this instant almost. There’s a lot of neutrinos, but they don’t interact with – with us, and so we’re completely – mostly unaware of it.

[slide animates on a new bullet point – From now on, the temperature of the universe will refer to the temperature of the photons only (radiation). Photons, neutrinos, and antineutrinos are the dominate components of the universe]

Okay, so let’s see. So, from now on, when we talk about the temperature of the universe, we’re going to be talking about the temperature of the – of the gassiest components of the universe.

[slide animates on a new bullet point – Temperature is too high for atomic nuclei to be stable]

And temperature is still too high for atomic nuclei to be stable –

[slide animates on a new bullet point – Neutrons are still being converted to protons plus electrons and the present values are N(n) approximately 17 percent and N(p) approximately 83 percent]

– and neutrons are still decreasing with time.

[new slide titled, Time equals three minutes and forty-five seconds, the Period of Nucleosynthesis, and featuring a bullet point that the temperature is now approximately 9 times 10 to the eighth Kelvin]

So, we go to three minutes and 45 seconds. Temperature is now about nine times 10 to the 8th Kelvin. Still very hot.

[slide animates on a new bullet point – Deuterium becomes stable (that is its binding energy is able to withstand bombardment of protons, neutrons, and photons, permitting heavier nuclei (3H, 3He, Li & Be) to be formed. This rapidly incorporates free neutrons into atomic nuclei where they are safe from further decay, fixing the N/P ratio, which determines how much He and D can be formed. He/H 2x (N/P) 26-28% by mass]

Deuterium can become stable, and so, but unfortunately its binding energy is not able to withstand the bombardment at these high temperatures. And so, it doesn’t actually form at this – it could with as far as the temperature is concerned, but – I’m mean, sorry. It couldn’t because the temperature is still too high, the bombardment energies are too high.

[new slide titled, The Period of Nucleosynthesis, featuring a graph with Temperature in Kelvin on the x-axis and Mass on the y-axis, showing mass increasing as temperature increases until about one time ten to the nineth Kelvin where the masses of elements level off and become steady]

Okay, if you just plot as temperature versus mass, there’s this short period where all these – these light elements can be formed. And this kind of started out with protons and neutrons, and at – at this point in time, our temperature of about 10 to the 9th Kelvin, you know, all these things get to be made –

[new slide titled, Time equals 34 minutes and 40 seconds, The Photon Dominated Universe begins with a bullet point that the temperature is now approximately three times ten to the eighth Kelvin]

– and it’s over a very short time. Okay –

[slide animates on a new bullet point – Matter density is now too low to support nuclear processing, even though the temperature is high enough. All elements capable of being produced in the Big Bang have been made]

– 34 minutes and 40 seconds. Proton dominated – we have a proton dominated universe. Temperature is about three times 10 to the 8th. Matter density is now too low to support nuclear processing, even though the temperature is high enough. All elements capable of being produced in the Big Bang have been made.

[slide animates on a new bullet point – e positive and e negative are completely annihilated, leaving only a small residual of about one in ten to the nineth of the original number. The same is also true of neutrons and protons. All that is left is the small matter-antimatter imbalance]

The electrons and protons or electrons and positrons have been completely annihilated, leaving only a small residual of about one in a billion of the original number. The same as protons and neutrons. And that’s – then this – this tiny residual is why we’re here because there was this imbalance, and, as I say, I – I – we don’t really quite understand why that imbalance was there. Okay.

[slide animates on a new bullet point – The main constituents of the universe are photons, neutrinos, and antineutrinos. Density is approximately 9.9 grams per cubic centimeter squared and 31% is in neutrinos and 69% is in photons – a photon dominated universe!]

So, the consistency of the universe now is photons, neutrinos, antineutrinos. Density is about 10 grams per cubic centimeter. And –

[Edward Churchwell, on-camera]

– of this density, about 31% is in neutrinos and 69% is in photons. Okay, I want to kind of move on because I’m going to run out of time here. So, the doubling times now of the universe is up to about one and a half hours. Every one and half hours, the universe gets – doubles in size. Okay, seven times 10 to the five years.

[slide titled, Time equals 7 times 10 to the fifth years, the Dark Ages begins]

So, we’re jumping along. This is called the Dark Ages.

[slide animates on the bullet point that the temperatures is about 3000 degrees Kelvin]

Temperature now is reduced to about 3,000 Kelvin.

[slide animates on a new bullet point – At this temperature, the universe is cool enough for electrons to be bound to atomic nuclei to form stable neutral atoms and the universe becomes transparent to photons. This is the beginning of what is called the dark ages]

At this temperature, the universe is cool enough for electrons to be bound to atomic – to – to nuclei. And so, you have atoms formed. And once that’s true, then suddenly the universe becomes pretty much transparent. Okay? And so, this is a period when before star formation has really started, and there’s gas that becomes atomic. And, as a consequence, you have a – a dark universe. Okay, and so it’s called the Dark Ages.

[slide animates on a new bullet point – This is the change from a radiation to a matter dominated universe]

And this is a change from a radiation to a matter dominated universe, okay. Up until this time, the –

[Edward Churchwell, on-camera]

– the largest number of particles in the universe was photons and neutrinos. Since neutrinos don’t interact, we’re not going to say much more about them here. But, at this point, when you start to form atoms, matter takes over as the dominant form of – of particles in the universe. Okay, now we go to about a billion years. This is a period of star galaxy formation. Temperature has now dropped to about 70 Kelvin or about minus 200 degrees Centigrade. Gravity has gathered matter into filaments and massive clouds that give rise to the formation of the first stars and galaxies. And that’s at an age of about a billion years.

At this time, it becomes possible for carbon-based life forms to appear in the universe for the first time. About nine billion years ago, the solar system is formed in the Milky Way. And from about one to nine billion years, it’s possible that life could have emerged –

[slide titled, approximately nine billion years, and featuring two bullet points. The first – the solar system is formed in the Milky Way; the second – From 1 to 9 billion years it is possible that life could have emerged many places in the universe. Life on Earth arose at age approximately 10 billion years (approximately 500 million years after Earths formation). But this would have been prokaryotic life. Eukaryotic life forms (reptiles, mammals, birds, etc.) are though to have formed about 2 billion years later. Humans are believed to have appeared more than a billion years later (at age 13 billion years)]

– many places in the universe. Life on Earth arose at the age of about 10 billion years. About 500 million years after Earth’s formation. But this would have undoubtedly been prokaryotic lifeforms, not – not the lifeforms like were – we’re made out of, but more primitive lifeforms. The eukaryotic lifeforms didn’t appear for another two billion years or so.

[new slide titled, Current Age of the Universe, featuring two bullet points – the first – approximately 13.8 billion years; the second – the temperature is approximately 2.73 degrees Kelvin or minus 270 degrees Celsius]

Okay, current age of the universe is about 13.8 billion years, and the current temperature of the universe is 2.73 degrees.

[new slide with two more bullet points – the first – Temperature, age, and mean separation of cosmological tracers are related; the second – Temperature r decreases as 1/t1/2. Size increases with age but the rate depends on what tracer dominates the universe (epoch)]

And I don’t think I need –

[new slide titled, Evolution: radiation, mass, and temperature dependence on time, featuring a graph with time on the x-axis and temperature on the y-axis showing temperature decreasing over time]

– to say any more about this. This is the – what I was trying to talk about a while ago, though. Time versus temperature. And the density of radiation is going down at a steeper rate than the density of matter. And you see there’s a crossing point here, and that’s when you convert from a radiation-dominated universe to a matter-dominated universe. The big T here is just the change in temperature with time. And that’s determined entirely by the cooling due to expansion.

[new slide titled, The Friedmann Equation, Alexander Freidman, 1922, used Einsteins GR Field Equations]

Okay, now I’m going to do something that I should never do in a general audience. I’m going to write an equation down; it’s a fairly complicated one. But I’m going to tell you – try to explain what’s going on here.

[slide animates on the equation, Ht2/H02 = Omega r,0 divided by a to the fourth plus Omega m,0 divided by a cubed plus Omega lambda plus one minus Omega 0 divided by a squared. Omega 0 is the present total energy density]

This is a – a relationship that was derived by Alexander Friedmann, who was a Russian meteorologist who taught himself relativity theory and then started working on cosmology and wrote this very beautiful equation. It’s called the Friedmann equation and basically what it is is what’s call an equation of motion of the – of the universe. And so, this H-naught we’ve already talked about, and the H sub-T squared is the Hubble constant at any other given time in any other time. Remember, it is a function of time. The rate – the universe is not expanding uniformly. It’s – its slowing down, or at least we thought this up to a few years ago.

Okay, so that is determined by the sum of the omegas. Omega R, omega radiation, omega mass, and the dark energy, omega lambda. But it has this dependence on this A factor. The A factor is what’s called the scale factor of the universe. And basically, all it is is the average separation of – on a cosmic scale of globular – of clusters of galaxies, okay? So, the average separation of – of clusters of galaxies, of course, is getting larger as time goes along. So, that’s a function of time. It’s just basically just determined by the expansion of the universe. And you can see that the radiation term here goes as this scale factor to the fourth power in the denominator. And so, it decreases very fast in its importance in the motion of the universe.

Similarly, for mass, omega M here is divided by A to the third power. And so, it decreases very fast too. Not as fast as radiation, but as we saw, we – we – the universe did convert from a matter-dominated to a – to a radiation-dominated to a matter-dominated universe. That’s because of that dependence on A here. Omega lambda doesn’t –

[Edward Churchwell, on-camera]

– depend on A at all. As a consequence, at some point it becomes the dominant factor in the evolution of the universe. And this one minus omega naught – omega naught is the sum of the other three here. So, we’re making the assumption that we know everything that’s in the universe is either photons, matter, or dark energy.

[return to The Friedmann Equation Slide]

Well, we – we think we’ve determined what these three quantities here are. And so, omega, we believe, is – is one. And so, that’s the term that determines the geometry of the universe. And if omega is one, that last term goes to zero, and that tells us that we live in a Euclidean universe. The universe is flat in that sense.

Okay, I’m going to –

[slide animates on a new bullet point – a is the cosmic scale factor which is the mean separation of particles (galaxies) in the Universe. This changes from very, very small values to very, very large values as the universe ages and explains why the universe was dominated by radiation in its early history, then matter, and finally dark matter]

– this cosmic scale factor, I’ve already said –

[slide animates on a new bullet point – The Friedmann equation describes the dependence of the components of the universe and its geometry on time and the cosmic scale factor a]

– what that is.

[slide animates on a new bullet point – The Friedmann equation can also be written as a sub-t equals H-naught times Omega r-naught divided by a squared plus Omega m-naught divided by a plus Omega lambda times a squared to the power]

You can convert this to a much simpler kind of expression. Whereas this is the rate of – of increase of the scale factor. So, just, you know, like, it’s a speed. And you can see it depends, when you write it this way, you still see the dependence on A and that and this – when you write it this way, you see that omega lambda is multiplied by A squared. The others have A in the denominator, so they go down very – rather fast. Okay.

Well, anyway, that tells us, though, that omega lambda is – is going to become, and in fact is now, the dominant – the dominant player in the evolution of the universe.

[new slice titled, The Future Evolution of the Universe, featuring three bullet points – To answer this question, we have to find a tracer that can be observed to very large distances and that is also a standard candle – A standard candle is an object that has an intrinsic and well-known brightness. A lighthouse provides a standard candle for ships – Type la supernovae fill these requirements. They result from a White Dwarf in a binary star system with a red dwarf star that transfers mass to the White Dwarf. When the mass added to the White Dwarf becomes large enough to cause the White Dwarf to shrink, causing its central temperature to become large enough to explosively ignite the C, N, O, and heavier elements, the White, Dwarf becomes a Type la supernova]

Okay, I’m going to kind of skip some of this –

[new slide titled, Brightness of Type la Supernovae vs. Time, featuring a graph with L/L0 on the y-axis and time on the x-axis]

– because –

[new slide titled, Current Standard Model Values from the Plank Mission and Type la SNRs, featuring five bullet points. 1)Temperature-naught equals 2.7525 plus or minus .0001 Kelvin; H-naught equals 67.3 kilometers per second per Mpc. T-naught equals 13.8 billion years, q-naught is less than zero = accelerating universe. 2) Omega r-naught equals 8.4 times 10 to the minus 5 Radiation; Omega lambda-naught equals 5 times ten to the minus 5 = photons; Omega v-naught = 3.4 times 10 to the minus 5 = neutrinos. 3) Omega m-naught = .315 Matter; Omega bary-naught = .049 Baryons; Omega dm-naught =.266 Dark Matter. 4) Omega lambda-naught = .685 = Dark Energy 5) 1-Omega-naught = 0 = a flat universe (Euclidian geometry) with a note that this assumes that we have accounted for all major components of the universe]

– I don’t want to – I want to tell you what the present, what’s called the standard model of the universe is today from all the measurements that weve –

[Edward Churchwell, on-camera]

– the best measurements we’ve been able to make.

The current temperature of the universe is 2.725 plus –

[return to the Current Standard Model Values from the Planck Mission and Type la SNRs slide]

– or minus 0.001 Kelvin. H-naught is 67.3 kilometers per second for every three billion – three million light years distance, increase in distance. And the age of the universe, as I say, is about 13.8 billion years. And Q-naught is less than zero, and we’re going to see that that tells us that the universe is not slowing down. It’s actually accelerating in its expansion because lambda, the omega lambda is the dominant term now –

[Edward Churchwell, on-camera]

– and it behaves, in a sense, like antigravity, if you wish.

So, the amount – the – the radiation term here total is only eight times 10 to the minus five. This is a ratio remember. It’s the – its the fraction of energy density that’s in the form of radiation. And that is almost – its – it’s not much. 10 to the minus five, 10 to the minus four. In photons, its five times 10 to the minus five. And in neutrinos, it’s three times 10 to the minus five. The sum gives you 8.4. So, photons are not contributing very much at the moment to the evolution of the universe.

Matter has a value of about 0.315, and of that, baryons only contribute about 5%. Baryons are the stuff we’re made out of. It’s the stuff stars are made out of. Galaxies are made out of. You know, it’s us. And so, most of what we know about the universe –

[return to the Current Standard Model Values slide]

– is only determined by about 5% of the contents of the universe, okay? So, that right away should bring – bring some red lights to your – to your head.

The dark matter, omega dark matter, is also part of the matter. And it makes up almost 27%. So, the stuff that we call dark matter, we don’t know what it is, is 27% of the contents of the universe.

Omega lambda, the dark energy, is almost 70% of the universe. And since –

[Edward Churchwell, on-camera]

– if you sum these up, you get about one. So that means omega naught is – is one, and, consequently, the – the geometry of the universe is what we call a flat universe or a Euclidean geometry.

Now, be aware that we’re assuming that we’ve included all the content. That we know what all the contents of the universe is. And that ain’t necessarily so, possibly. It’s an uncertainty.

[slide titled, Type la S.N.R.s = An Accelerating Universe, featuring a graph with Omega-mass on the x-axis and Omega-lambda on the y-axis, showing a sort of half oval line pattern between Omega-1 lambda and Omega-naught lambda between Omega-mass zero and 2.0]

I’m just going to show you this. It’s kind of complicated.

[new slide titled, Graphic of Universes Evolution, featuring two graphs. At the top, is a graph with a variable wave form with gravitational values at the top and density values at the bottom. At the bottom is a graph of radiation values of the universe from inception to the creation of galaxies]

I’m going to just zip over it. So, I’m going to give you a thumbnail picture of the evolution of the universe now, starting at the Big Bang right here, which we – I didn’t say anything about because we don’t know anything about it. But we had this inflation period where the universe just explosively expanded by maybe 25 orders of magnitude in a tiny, tiny fraction of a second. And then we had the formation of – of protons. And then we had nuclear fusion at these various ages. Time is running along on the bottom here. And think of the size of the universe, if you wish, as the – the height. Okay, well, in any case, this is kind of a thumbnail picture of what we’ve been talking about.

[new slide titled, What about the Future, Four Possibilities, featuring four illustrations, the first is titled – the re-collapsing universe – shaped like an oval where in the future it will collapse in on itself and disappear; the second is titled – the critical universe – shaped like a bell where the universe will keep going at its present rate; the third is titled – the coasting universe – shaped as a V where the universe with expand at a slightly faster rate; the fourth is titled – the accelerating universe – shaped like a funnel with the universe expanding exponentially over time]

And let’s just now ask about the future. And there were four possibilities. If you think about the present as right here. When I was a grad student, we used to argue about whether the universe, although we knew that it’s expanding now, its own gravity may have been enough density of matter, it was possibly large enough that it could actually collapse on itself and go back into a – a singularity. Other possibility what’s called a critical universe, where it’s expanding outward as it goes, but it’s never – there’s never large enough density of matter to actually cause it either to come completely to a stop or certainly not fall on itself. And then a coasting universe where it just continues out uniformly expanding. And the one that now seems to be the best assumption is the accelerating universe here because dark energy has taken over and the universe is actually accelerating at today’s, at the present time. So, that’s the one that the standard model now predicts. But –

[new slide titled, Many Remaining Questions]

– there are many questions.

[slide animates on the question – How much might we be missing by only observing 5% of the contents of the Universe?]

How much might we be missing by only observing 5% of the – of the contents of the universe?

[slide animates on the question – Might there be other undiscovered constituents? That is, is Omega-naught not equal to 1 as the current model implies]

Might there be other undiscovered constituents? That is, is omega-naught not equal to one as a current standard model would predict?

[slide animates on the question – what is dark matter?]

What is dark matter? We don’t really know what that is.

[slide animates on the question – what is dark energy? Why does dark energy behave like anti-gravity? Why is dark energy independent of the universal scale factor, whereas all other constituents are dependent on it?]

What is dark energy? We also, you know – why does dark energy behave like anti-gravity? Why is dark energy independent of the universal scale factor? As everything else, all other constituents in the universe depends –

[Edward Churchwell, on-camera]

– on the scale factor but lambda – omega-lambda does not. And that has to go back to the mathematics of when Einstein introduced this lambda term. And he introduced it because the Friedmann equation was un- unstable. So, anyway, he thought it was his – his greatest error –

[return to the Many Remaining Questions slide]

– but it almost certainly was not.

[slide animates on the question – If neutrinos have mass, as experiments seem to imply, why have anti-neutrinos not disappeared from the Universe?]

If neutrinos have mass, as experiments seem to imply, why have anti-neutrinos not disappeared from the universe? Another good question.

[slide animates on the question – What triggered the inflation of the Universe and what stopped it? We know essentially nothing about the Planck period and the laws of nature that govern the interactions of particles at this stage and the force(s) that govern their interactions]

What triggered the inflation of the universe, and – and what stopped it? We know basically essentially nothing about the Planck period and the laws of nature that govern the interaction of particles at this very earliest stage, and we probably are not going to –

[Edward Churchwell, on-camera]

– know much more about it until we understand the quantum nature of gravity. Gravity is the one force that we still, believe it or not, it sounds simple but it ain’t, believe me. It’s a very complicated issue, and we – we don’t know how it would behave in a quantum sense. And so, that’s still a big open question. Fundamentally, why is the universe expanding? Okay, well, I should stop there. I think I’ve taken your – plenty of time. So, thank you very much. I will be glad to take questions.

[applause]