II Statistical Physics - Quantum gases

3Quantum gases

II Statistical Physics

3.7 Fermions

We now move on to study fermions. One major application of this is the study

of electrons in a metal, which we can model as gases. But for the moment, we

can work in complete generality, and suppose we again have a non-interacting

fermion gas. Since they are fermions, they obey the Pauli exclusion principle.

Two fermions cannot be in the same state. Mathematically, this means the

occupation numbers

cannot be arbitrary — they are either 0 or 1. As in

the case of bosons, lets study these fermions using the grand canonical ensemble.

The partition function is

Z =

where

−βn

−µ)

= 1 + e

−β(E

−µ)

This is just like the bosonic case, but we only have two terms in the sum. We

then find that

log Z =

log(1 + e

−β(E

−µ)

Then we have

i = −

∂

∂E

log Z =

β(E

−µ)

+ 1

This is called the Fermi-Dirac distribution. The difference between this and the

bosonic case is that we have a difference in sign!

We can get the total number of particles in the usual way,

hNi =

∂

∂β

log Z =

β(E

−µ)

+ 1

hri.

As before, we will stop drawing the h·i.

There is another difference between this and the bosonic case. In the bosonic

case, we required

µ <

0 for the sum

to converge. However, here the sum

always converges, are there are just two terms. So there is no restriction on µ.

We assume that our particles are non-relativistic, so the energy is

E =

Another thing to notice is that our particles have a spin

s ∈ Z +

This gives a degeneracy of

= 2

+ 1. This means the density of states has to

be multiplied by g

. We have

g(E) =

4π





3/2

1/2

We can now write down the total number of particles in terms of the density of

states. We have

N =

∞

dE g(E)

−1

βE

+ 1

where again

z = e

βµ

Similarly, the energy is given by

E =

∞

dE Eg(E)

−1

βE

+ 1

Again, we have the density of states

pV = kT log Z = kT

∞

g(E) log(1 + ze

−βE

Just as for bosons, we can integrate this by parts, and we obtain

pV =

We now look at the high temperature limit. In the case of boson, we identified

this with a small

, and we can do the same here. This is done on example sheet

3, and we find that to leading order, we have

z ≈ λ

exactly as for bosons. This is exactly the same as in bosons, and again as in the

example sheet, we find

pV = N kT



1 +

√

+ ···



We see the sign there is positive. So Fermi statistics increases pressure. This is

the opposite of what we saw for bosons, which makes sense.

We now go to the opposite extreme, with zero temperature. Then

∞

β(E−µ)

+ 1

(

1 E < µ

0 E > µ

Note that here the

is evaluated at 0, as

depends on

. Clearly the value of

µ at T = 0 is important.

Definition (Fermi energy). The Fermi energy is

= µ(T = 0) = lim

T →0

µ(T, V, N ).

Here we see that at

= 0, all states with

E ≤ E

are occupied, and states

with higher energy are not occupied. Now it is easy to work out what this Fermi

energy is, because we know the total number of particles is N. We have

N =

dE g(E) =

6π





3/2

Inverting this, we find that



6π



2/3

We note that this is very much like the critical temperature of the Bose–Einstein

distribution. We can define the characteristic temperature scale

Now whenever we talk about whether the temperature is high or low, we mean

relative to this scale.

How high is this temperature? We saw that for bosons, this temperature is

very low.

Example. For

He, the Fermi temperature is

∼ 4.3 K.

While this is very low compared to everyday temperatures, it is pretty high

compared to our previous T

For electrons in a metal, we have

∼ 10

. This is very high! Thus, when

studying metals, we cannot take the classical limit.

For electrons in a white dwarf stars, we have

∼ 10

! One reason this is

big is that

is the electron mass, which is very tiny. So it follows that

very large.

From the Fermi energy, we can define a momentum scale

= (2mE

)

1/2

While we will not be using this picture, one way to understand this is that in

momentum space, at

= 0, all states with

|k| ≤ k

are occupied, and these

states are known as a Fermi sea. States at the boundary of this region are known

as Fermi surface.

Let’s continue working at zero temperature and work out the equations of

state. We have

E =

dE Eg(E) =

This is the total energy at zero temperature. We saw earlier that

pV =

E =

This shows the pressure is non-zero, even at zero temperature! This is called the

degeneracy pressure. This is pressure coming from the Fermi statistics, which is

completely unlike the bosons.

Now let’s see what happens when we are at low, but non-zero temperature,

say

T  T

. While we call this low temperature, since we saw that

can be

pretty high, this “low temperature” might still be reasonably high by everyday

standards.

In this case, we have

≈

 1.

Therefore we find that z  1.

Our goal now is to compute the heat capacity. We begin by trying to plot

the Fermi-Dirac distribution at different temperature scales based on heuristic

arguments. At T = 0, it is simply given by

n(E)

When we increase the temperature a bit, we would expect something that looks

n(E)

Since

is small compared to

, that transition part is very small. Only states

within kT of E

are affected by T .

We can now give a heuristic argument for what the capacity should look like.

The number of states in that range is

∼ g

(

)

. Relative to where they are in

= 0, each of these is contributing an excess energy of order

. Therefore the

total energy is

E ≈ E|

T =0

+ g(E

)(kT )

So we expect the heat capacity to be

∼ k

g(E

)T ∼ kN

This rather crude argument suggests that this heat capacity is linear in tempera-

ture. We shall now show that this is indeed correct by direct calculation. It is

the same kind of computations as we did for bosons, except it is slightly more

complicated. We have

4π





2/3

∞

1/2

−1

βE

+ 1

Similarly, we find

4π





3/2

∞

dE E

3/2

−1

βE

+ 1

We again do a substitution x = βE, and then we find

3/2

(z)

kT f

5/2

(z),

where we similarly have

(z) =

Γ(n)

∞

n−1

−1

+ 1

Recall that we have





√

, Γ





√

We know that

z 

1, and we want to expand

. This is called the

Sommerfeld expansion. We have

Γ(n)f

(z) =

βµ

n−1

−1

+ 1

∞

βµ

n−1

−1

+ 1

| {z }

(2)

We can rewrite the first term as

βµ

dx x

n−1



1 −

1 + ze

−x



(log z)

−

βµ

n−1

1 + ze

−x

| {z }

(1)

using the fact that z = e

βµ

We now see that the two remaining integrals (1) and (2) look very similar.

We are going to make a change of variables to make them look more-or-less the

same. We let

= βµ − x

in (1), and

= x − βµ

in (2). Then we find that

Γ(n)f

(z) =

(log z)

−

βµ

dη

(βµ − η

)

n−1

1 + e

∞

dη

(βµ + η

)

n−1

+ 1

Since

βµ 

1, we may approximate the first integral in this expression by

∞

The error is ∼ e

−βµ

∼

 1, so it is fine.

Calling η

= η

= η, we get

Γ(n)f

(z) =

(log z)

∞

dη

(βµ + η)

n−1

− (βµ − η)

n−1

1 + e

Let’s now look at the numerator in this expression here. We have

(βµ + η)

n−1

− (βµ − η)

n−1

= (βµ)

n−1



1 +

βµ



n−1

−



1 −

βµ



n−1

= (βµ)

n−1



1 +

(n − 1)η

βµ

+ ···





1 −

(n − 1)η

βµ

+ ···



= 2(n − 1)(βµ)

n−2



1 + O



βµ



We just plug this into our integral, and we find

Γ(n)f

(z) =

(log z)

+ 2(n − 1)(log z)

n−2

∞

dη

η+1

+ ··· .

We have to be slightly careful. Our expansion was in

βµ

, and while

βµ

is large,

can get arbitrarily large in the integral. Fortunately, for large

, the integrand

is exponentially suppressed.

We can compute

∞

dη

+ 1

∞

dη

ηe

−η

1 + e

−η

∞

dη η

∞

m=1

−mη

(−1)

m+1

∞

m=1

(−1)

m+1

Note that it was valid to swap the sum and integral because the sum is absolutely

convergent.

So we find that

(z) =

(log z)

Γ(n + 1)



1 +

n(n − 1)

(log z)

+ ···



Hence we find that

6π

(2mµ)

3/2

1 +





+ ···

. (∗)

Recall that the reason we did this is that we didn’t like

. We want to express

everything in terms of

instead. It is an exercise on the example sheet to invert

this, and find

µ = E

1 −





+ ···

. (†)

We can plug this into the epxression of

and obtain

10π

(2m)

3/2

5/2

1 +

5π





+ ···

Dividing by (∗) and using (†), we find

1 +

5π





+ ···

We can then differentiate this with respect to

, and find the heat capacity to



∂E

∂T



V,N

+ ··· ,

and this is valid for T  T

Heat capacity of metals

In a metal, we have a lattice of positive charges, and electrons that are free to

move throughout the lattice. We can model them as a Fermi ideal gas. As we

mentioned previously, we have

∼ 10

So we usually have

T  T

. In particular, this means any classical treatment of

the electrons would be totally rubbish! We have the lattice as well, and we have

to think about the phonons. The Debye temperature is

∼ 10

Now if

 T  T

then, assuming an equal number of electrons and atoms, we have

≈

| {z }

electrons

+3Nk ≈ 3N k.

This agrees with the Dulong–Petit result.

But if

T  T

, T

, then we need to use the low energy formula for the

phonons, and we get

≈

12π





In this case, we cannot just throw the terms away. We can ask ourselves when

the contributions of the terms are comparable. When are their contributions

comparable. This is not hard to figure out. We just have to equate them and

solve for T , and we find that this happens at

24π

If we plug in actual values, we find that this is just a few Kelvins.

While this agrees with experiments very well, there are some flaws. Most

notably, we neglected interactions with the lattice. We need to correct for a

periodic potential, and is not too difficult. This is addressed in the AQM course.

Even worse, we are ignoring the interaction between electrons! It is a puzzle

why our model works despite ignoring these. This is explained by the Landau–

Fermi liquid theory.

White dwarf stars

Another important application of this is to astrophysics. A star is a big ball

of hot gas that is not collapsing under gravity. Why does it not collapse under

gravity?

In our sun, there are nuclear reactions happening in the core of the star,

and these produce pressure that resists gravitational collapse. But this process

eventually ends. As the star runs out of fuel, the star will shrink. As this

happens, the density of the star increases, and thus the Fermi energy

for

electrons increase. Eventually this becomes so big that it exceeds the energy

required to ionize the atoms in the star. This happens at around

∼ 10

At this stage, then, the electrons are no longer bound to the atoms, and all

electrons can move through the star. We can therefore attempt to model these

electrons as a low temperature Fermi gas (because T

is now very high).

One important feature of Fermi gases is that they have pressure even at zero

temperature. In a white dwarf star, the degeneracy pressure of this gas supports

the star against gravitational collapse, even if the star cools down to

= 0! A

particular important example of a star we think will end this way is our own

sun.

Let’s try and understand some properties of these stars. We construct a very

crude model of a white dwarf star by treating it as a ball of constant density. In

addition to the kinds of energies we’ve talked about, there is also gravitational

energy. This is given by

grav

= −

3GM

This is just a statement about any star of constant density. To see this, we ask

ourselves what energy it takes to destroy the star by moving shells of mass

δM

off to infinity. The work done is just

δW = −PE =

δM.

So we have

R(M)

(M/(

πρ))

1/3

where

is the fixed density. Solving this, and using the expression of

in terms

of M gives

−E

grav

= W =

3GM

Let’s see what stuff we have got in our star. We have

electrons, hence

protons. This gives

αN

neutrons for some

α ∼

1. In the third example sheet,

we argue that we can ignore the effect of the protons and neutrons, for reasons

related to their mass.

Thus, we are left with electrons, and the total energy is

total

= E

grav

+ E

kinetic

We can then ask ourselves what radius minimizes the energy. We find that

R ∼ M

−1/3

∼ N

−1/3

In other words, the more massive the star is, the smaller it is! This is pretty

unusual, because we usually expect more massive things to be bigger.

Now we can ask ourselves — how massive can the star be? As the mass goes

up, the radius increases, and the density goes up. So

goes up. For more and

more massive stars, the Fermi energy is larger and larger. Eventually, this

becomes comparable to the rest mass of the electron, and relativity becomes

important.

Consider E

 mc

. Then we have E  mc

, and we can expand

g(E) ≈



−

+ ···



where we plugged in g

= 2. We can work out the number of electrons to be

N ≈

dE g(E) =



−



Here we are assuming that we are at “low” energies, because the Fermi tempera-

ture is very high.

Recall this is the equation that gives us what the Fermi energy is. We can

invert this to get

= ~c



3π



1/3

+ O





1/3

Similarly, we find

kinetic

≈

dE Eg(E)



−

+ ···



4π

(3π

4/3

1/3

+ O





−2/3

We now note that

πR

, and

m  m

≈ m

, where

m, m

, m

are the

masses of an electron, proton and neutron respectively. Then we find that

total

3~c



9πM

4(α + 1)



1/3

−

3GM

+ O(R).

This is the total energy once the electrons are very energetic. We can plot what

this looks like. If the coefficient of

is positive, then we get something that

looks like this:

There is a stable minimum in there, and this looks good. However, if the

coefficient is negative, then we go into trouble. It looks like this:

This is bad! This collapses to small

. In other words, the degeneracy of the

electrons is not enough to withstand the gravitational collapse.

For stability, we need the coefficient of

to be positive. This gives us a

maximum possible mass

∼





3/2

This is called the Chandrasekhar limit.

Should we believe this? The above calculation neglects the pressure. We

assumed the density is constant, but if this were true, the pressure is also

constant, but this is nonsense, because a constant pressure means there is no

force!

A more careful treatment shows that

≈ 1.4M

sun

What if we are bigger than this? In the case of a very very very dense star, we

get a neutron star, where it is the neutrons that support the star. In this case,

we have protons combining with electrons to get neutrons. This would involve

actually doing some general relativity, which we will not do this.

This neutron star also has a maximum mass, which is now

max

≈ 25M

sun

If we are bigger than this, then we have to either explode and lose a lot of mass,

or become a black hole.