Lesson learning outcomes:

By the end of this lesson, you should be able to:

1.   Describe the structure of the Sun, including the energy fusion process in the core and its different layers.

2.   Describe how the Sun regulates energy via its thermostat .

3.   Recognise that the Sun has a magnetic field and use this to describe the presence of sunspots.

4.   Explain how solar flares and coronal mass ejections occur.

7.1 Introduction to the Sun

The Sun is the most important thing that makes life on Earth possible. In this lesson, we are going to look closer at our host star.

About the photo: NASA’s Solar Dynamics Observatory captured this x-ray image of our Sun, with a solar flare in

progress (bright spot at far leftjust above centre). Colours correspond to the strength of the x-ray emission, with the

brightest emission in blue and purple.

We have learned all about the delicate balance of gravity and internal gas pressure that is occurring inside stars, and our Sun is no different. The Sun is a main sequence star (we learned about main sequence stars in Lesson 6) and therefore fuses hydrogen atoms via the proton-proton (p-p) chain in its core. This is what we call gravitational equilibrium.

There is another kind of equilibrium that gravitational equilibrium is related to, and this is energy balance. The energy produced in the Sun’s core due to nuclear fusion balances the energy that radiates from the Sun’s surface. We are going to learn more about energy transfer in this lesson.

Thanks to the fact that the Sun is so close to us, we can learn a lot more about it than we can other stars, including its structure. We will talk more about this later, but let’s start beyond the surface. As spoke about last lesson, a solar wind is often produced by stars, and this is the same for the Sun. In fact, the solar wind, which is a stream of particles blown outward from the Sun at all times, helps shape the magnetosphere of planets and the shape of comet tails, but we will learn about these in later lessons.

Heading further in, we come across the corona. The corona is very low-density gas and is the uppermost layer of the Sun’s atmosphere. The corona emits most of the Sun’s x-rays, and this is unsurprising given that it is about 1 million K: much, much higher than the surface of the Sun. The reason for this extremely high temperature is still an active field of research.

As we  head  closer to  the  surface,  the  temperature  drops  from  1,000,000  K  in  the  corona  to  10,000  K  in  the chromosphere. Most of the UV light radiated from the Sun is from this region. The visible surface of the Sun though is further down. It is called the photosphere and is the lowest layer of the Sun’s atmosphere. It is only around 6,000 K. Spectroscopy of the surface confirms that the Sun is predominantly hydrogen and helium, as we would expect. Also on the surface are sunspots, which can often be larger than the size of the Earth. The radius of the Sun is more than 100 times the radius of Earth.

The mass of the Sun can be calculated using Newton’s version of Kepler’s third law, which we learned back in Lesson 2. Its mass is about 2 x 1030 kg, which is just under 1000 times the mass of all of the planets in the solar system combined (and about 300,000 times the mass of the Earth).

The Sun rotates; how do you think we know this? Well, several ways. Recall our discussion on the Doppler effect (also Lesson 2); well, we can use the Doppler effect for determining how fast each side of the Sun is either moving towards or away from us. Additionally, we can also track the movement of sunspots across the surface. The Sun’s rotation is not like a solid ball; instead, the equator completes one rotation in about 25 days, whereas the poles take closer to 30 days to complete a rotation.

Heading deeper into the Sun beyond the surface, we pass through the Sun’s convection zone first, then radiative zone. We spoke about convection briefly last lesson. Convection is the rising and falling of large plumes of material and is a very efficient way of transporting energy. Convection tends to arise when there are two layers of gas, such as the Sun’s interior and the Sun’s surface, that have vastly different temperatures. In an attempt to even things out, convection starts up as a way to try and transport material and energy. Convection near the surface is what gives the surface of the Sun its granulated look, which we will talk about soon. Further down, we hit the radiative zone just above the Sun’s core.

The Sun produces an incredible amount of energy in its core, ready to be radiated out into space. Each second, 3.8 x 10^26 Joules of energy are produced (in units of power, this is 3 x 10^26 W, and is the Sun’s luminosity, which we learned about in Lesson 6). This amount of energy would meet all of Earth’s energy needs for 500,000 years, if only we could just harness it. Of course, since this energy radiates out in all directions, only a small fraction reaches Earth.

7.2 Nuclear fusion in the Sun

We have spoken about fusion in previous lessons, so will not go into repetitive details here. Extremely high temperatures and pressures are required to fuse two atoms together. As we have spoken about previously, the Sun is a main- sequence star, which means that it is fusing hydrogen atoms together in its core via the proton-proton (p-p) chain. We spoke about the p-p chain when we were discussing main-sequence stars in Lesson 6. Also remember that while 4 hydrogen atoms fuse to create one helium nucleus, the mass of one helium nucleus is less than the combined mass of 4 protons, and that is because some mass is converted into energy, and this is what powers the Sun.

The Solar Thermostat

You might be wondering what is going to happen since the Sun is slowly losing mass, especially given that, up until this point, we have emphasised that everything in the way stars is structured is essentially a balancing act. It is very beneficial for life on Earth that the Sun’s luminosity remains relatively constant, as this is, we conjecture, the basis of life on Earth. The Sun’s luminosity is relatively constant thanks, in part, to the Sun’s thermostat (this thermostat is not unique to the Sun, however!).

You can probably guess how the Sun’s thermostat works based on everything we have discussed previously, but here it is for you:

-     If the Sun’s core temperature rises a little bit, then:

o  The hydrogen fusion rate in the core increases, which causes:

o  The thermal pressure in the core to increase, which means:

o  Thermal pressure will be greater than gravitational pressure, so:

o  The core will expand, cool down, and:

o  The fusion rate returns to normal.

-     If the Sun’s core temperature drops a little bit, then:

o  The hydrogen fusion rate in the core decreases, which causes:

o  The thermal pressure in the core to decrease, which means:

o  Thermal pressure will be less than gravitational pressure, so:

o  The core will contract, heat up, and:

-     The fusion rate returns to normal.

Figure 14.9 from The Cosmic Perspective. The solar thermostat.

The Sun is gradually becoming brighter over time. Why is this the case? As 4 hydrogen atoms become 1 helium nucleus, the total number of particles decreases. What this means for the core is that it tends to shrink over long periods of time, since there are physically less particles there. The core temperature rises during this shrinkage, which is a good thing, because since there are less particles to support the weight of the matter above the core, the fusion rate has to increase so that thermal pressure can be maintained. So, how much does the Sun’s luminosity actually rise? Since the Sun was born, astronomers estimate that the Sun’s brightness has increased by about 30%!

7.3 Energy and the Sun

It might surprise you to learn that energy generated in the core of the Sun does not instantaneously get radiated from the surface. Firstly, as we know, the energy produced in the Sun’s core is in the form of photons, so how do those photos make their way from the core to the surface? What do they encounter along the way that would make their path not a straightforward one?

You might be able to understand this better if you have ever tried to navigate crowded streets in a city. In this situation, you might want to make what would be a very straightforward, direct path from one place to another. However, as you bump into people and try to avoid particularly crowded areas and cars, your path will not be straight, even though you will eventually reach your destination. The analogy applies here (though not the specifics, obviously!).

When a photon starts its journey outward, it begins in the core, which is very dense. The atoms in the core are packed quite close together; so close in fact that a photon can only travel a fraction of a millimetre before it runs into another atom, or, rather, the electrons surrounding that atom. Sometimes, if the photon is the right energy, it will be absorbed by the electron and re-emitted, or it will simply be deflected. This incredibly slow migration outward from the centre takes place over hundreds of thousands of years (!) and is formally called radiative diffusion (you can think of diffusion as “spreading out”, like what happens when you put a dash of milk into tea, and it eventually mixes).

Now, it is important to recall the structure of the Sun that we learned about in 7.1. This radiative diffusion occurs in the radiative zone just above the Sun’s core. Above the radiative zone is the convection zone.

We alluded to convection in Lesson 6. Convection, as was stated earlier, is where hot material rises, and cool material falls in a circular-type motion. Material that is hotter is less dense because the particles travel faster and cause the material to expand. This means that material that is hotter than its surroundings will be less dense and will rise. As the hot material rises, it transfers energy to its surroundings and cools down. It becomes denser, causing it to sink again. This is quite an efficient form of energy transport, much more efficient than radiation.

You might be wondering; how do we know that convection is occurring in the Sun? Well, we know because the photosphere (the “surface” of the Sun that we can see because it is the layer in which photons escape into space) shows sign of convection. Hot gas is significantly brighter than cold gas, and the Sun’s visible surface shows pockets of rising gas surrounded by sinking gas (the pockets are called granules, and the surface’s appearance is technically called granulation).

Figure 14.11 from The Cosmic Perspective. The Sun’s photosphere churns with rising hot gas and falling cool gas as

a result of underlying convection.

7.4 The Solar Interior

So, how do we know what is going on in the Sun’s interior? There are three ways: computer models, earthquake-like oscillations in the Sun, and neutrinos emanating from the surface.

1.   Computer models

Astrophysicists can use the laws of physics and things that we already know about the Sun, such as the Sun’s current radius, surface temperature, luminosity, age, among other things, to try and predict what is happening below the surface. If computer models can accurately predict what we already know and obey the laws of physics, then we can place more confidence in those models.

2.   Oscillations

It is important to remind ourselves at this point that the Sun is a ball of gas, so all of the (rather) violent motions we were talking about before, i.e., convection, make the Sun change shape slightly.

The study of this phenomenon in the Sun is called helioseismology. The whole concept is very complicated; however, we will try to summarise it here.

To summarise the Sun’s oscillations, we will make reference to how earthquakes occur on Earth, which has similarities to oscillations in the Sun. When sudden rock movement occurs in the Earth, seismic waves carry energy through the Earth’s layers, and these waves propagate through different materials in different ways. In fact, it is so predictable how these seismic waves travel through different materials that we can accurately determine characteristics (such as density) of different layers of the Earth.

In the Sun, when violent convective motions occur, they cause a lot of waves. In fact, there are three different kinds of waves that helioseismologists measure or look for: acoustic, gravity, and surface gravity waves , which generate p, g, and f modes respectively. There are about 10 million p and f modes alone. Acoustic waves are a sound wave, so you can think of them sounding like an audience clapping! Also, these waves oscillate in frequencies that could be audible to the human ear, if there were matter in space for the waves to propagate through (you can listen to some Sun sounds here: http://solar-center.stanford.edu/singing/singing.html).  Convection  messes  up  the way  that  these  oscillations propagate, so it is difficult for astrophysicists to measure oscillations from the interior (below the convective region). Gravity waves, produced below the convective region, are very difficult to detect and measure.

Even though we could hear these vibrations, this is clearly not the way that they were discovered. They were actually first detected in the 1960’s by observing the Doppler shift of each part of the Sun’s surface and continuously measuring that over time. Computers detect patterns in the behaviour; hence we can learn about what is going on inside the Sun. Thankfully, the oscillations measured and the computer models agree very well, which indicates that we know quite well about what is going on inside the Sun.

Figure 14.12 from The Cosmic Perspective. This image shows vibrations on the Sun’s surface that have been

measured from Doppler shifts. Shades of orange show how quickly each spot on the Sun’s surface is moving toward

or away from us at a particular moment.

3.   Neutrinos

Remember how we learned about the p-p chain in Lesson 6? Here is again for your reference:

Figure 14.8 from The Cosmic Perspective. In the Sun, four hydrogen nuclei (protons) fuse into one helium-4 nucleus

by way of the proton-proton chain.

Well, you’ll notice that there are some neutrinos that have been produced in the process. In fact, neutrinos only interact with matter through the weak force and gravity (we spoke about these forces in Lesson 3; gravity is much weaker than the other forces unless at extremely high temperatures, and the weak force only acts over a very small distance). This means that neutrinos don’t generally interact with matter and will happily travel through materials, including your body! What this also means though is that neutrinos produced in the core of the Sun during the fusion of hydrogen atoms pass through the Sun’s layers undisturbed.

You might think, “Great, that’s an easy problem to solve; let’s just build some neutrino detectors and be done with it!” . That’s what early astrophysicists thought also, but there is an issue: since neutrinos do not readily interact with matter, how do we detect them? Well, since neutrinos do sometimes interact with matter, you have to build a very, very large detector. Neutrino detectors are typically built underground or under the Antarctic ice. Overlying rock and ice will block most particles, but neutrinos pass right through.

Neutrino detection began in the 1960’s. Initially, detectors found around 1/3 of the neutrinos expected from models of the Sun’s output, which became known as the solar neutrino problem. Around 40 years later, in the early 2000’s, the solution was found; neutrinos came in 3 types: electron, muon, and tau neutrinos . Early neutrino detectors could only detect electron neutrinos, but that was fine because that’s the only type of neutrino produced in fusion reactions (i .e., those in the core of the Sun). However, along their journey to us from the Sun’s core, neutrinos can change between the types, so by the time they have reached us, only about 1/3 of the neutrinos are still electron neutrinos. Modern neutrino detectors detect all three types and have confirmed our suspicions, hence the neutrino problem is solved!

Figure 14.15 from The Cosmic Perspective. This photo shows the inside of the main vessel of the Borexino neutrino

detector, located in a mine 1400 meters underground at Italy’s Gran Sasso National Lab.

7.5 The Sun’s magnetic field

One of the final interesting points that we are yet to touch upon for the Sun are its sunspots. Sunspots are part of a broader property of the Sun, which is its “weather”, or solar activity. Storms associated with solar activity or solar weather affect our lives on Earth, and you’ve probably read some articles about this.

We saw earlier that the Sun has granulation on its surface due to the rising and falling of gas during its convective motions. This is very small-scale activity. There are also much larger events that occur on the solar surface, such as solar flares and sunspots, and they share a common cause: the Sun’s magnetic field.

Sunspots

Let us firstly look at sunspots. Sunspots are incredible features on the solar surface. They are visibly darker than their surroundings because they are much cooler (about 4000 K compared to 5800 K).

Figure 14.16 from The Cosmic Perspective. Sunspots are regions of strong magnetic fields.

But how does this happen? How can sunspots be cooler than their surroundings and not quickly be warmed up by the incredibly hot plasma that surrounds it? The answer points to magnetic fields.

Figure b) above shows that sunspots are regions of very strong magnetic fields. We know this because of the “line splitting” that has occurred in the absorption spectra lines of the sunspot (the splitting is formally called the “Zeeman effect”, but you do not need to remember the name). What is a magnetic field? You’ll know a magnetic field if you have ever played with magnets; it is an invisible field that certain materials create, and it is responsible for the difficulty you face when trying to push two of the same magnetic poles together as well as the ease of the coming together of two opposite poles. You can see the shape of a magnetic field around a bar magnet below.

Figure: Magnetic field around abar magnet.

We represent a magnet by drawing magnetic field lines. The strength of a magnetic field decreases with distance, and just like a contour map, where lines are drawn closer together, the field is stronger.

While the magnetic field lines are representative only, there is an interesting phenomenon that accompanies a magnetic field: charged particles tend to move along the field lines, rather than cutting them (perpendicularly).

The solar atmosphere is very turbulent, and the magnetic fields can be twisted up and stretched like long elastic bands. In places where the magnetic field is tightly wound, the magnetic field lines are very close together and the field is very strong. Convection, which occurs in the surface layers of the Sun, is suppressed in this region because the magnetic field is so strong and charged particles are moving along it. Hot plasma cannot enter the region where the magnetic field is very strong, meaning that the temperature drops, and a sunspot is formed.

Individual sunspots can typically last for up to a few weeks and will disappear when the magnetic field weakens, and plasma can flow into the region again .

Sunspots typically occur in pairs, however. The magnetic field produced under the surface layers arcs up and around, creating a giant loop. Gas in the chromosphere and corona are trapped in the loop, and we call this a solar prominence. Solar prominences rise to more than 100,000 km above the surface and can last for weeks.

We have observed particularly intense bursts of light and energy from the Sun (UV and x-rays) that originates from the vicinity of sunspots. These intense bursts are called solar flares . It is theorised that they are created when magnetic field lines get tangled; this causes the field lines to “break” . Remember that particles travel along the field lines, so when they break, particles flow from the surface.

Solar wind

In regions where the magnetic field lines are broken and do not repair, particles are able to stream freely into space. This is what we call the solar wind and studying the constituent particles allows us to capture a sample of the Sun!

Coronal mass ejections

“Normal” solar flares and solar wind are not an issue for us on Earth, because the Earth’s magnetosphere protects us. Sometimes, however, these events are particularly intense. When this happens, a significant amount of high-energy charged particles is ejected from the Sun’s surface; these are called coronal mass ejections (coronal – from the corona – mass ejections – ejections of actual particles). You’ve probably seen news articles about CMEs. Given that they are highly energetic and charged, they affect electrical components in satellites, disrupt power supply on Earth, and disrupt radio communications.

It isn’t all bad news for CMEs though. They excite particles in Earth’s atmosphere and can lead to some particularly strong auroras.

Sunspot cycle

While there is a lot that is incredibly unpredictable about the Sun’s activity, there is quite a bit that is predictable. Since around 1900, astronomers have observed the number of sunspots on the Sun’s surface and noticed that there is a pattern of how the number of prominences, flares, CMEs, and sunspots changes over time. Such solar events are common when solar activity is at its maximum, and are less frequent when at its minimum. The time between maximums is around 11 years, but varies from 7 to 15 years.

At the start of a solar minimum, sunspots will form at mid-latitudes (30-40 degrees – can you work out where this corresponds to on Earth? In Australia, 30S is around Coffs Harbour, NSW down to 40S, which is around Launceston in Tasmania). As the solar cycle progresses, the sunspots will form closer to the equator.

Figure 14.23 Sunspot cycle since about 1900.