Initial mass function

The IMF of stars more massive than our sun was first quantified by Edwin Salpeter in 1955.[6] His work favoured an exponent of ${\displaystyle \alpha =2.35}$. This form of the IMF is called the Salpeter function or a Salpeter IMF. It shows that the number of stars in each mass range decreases rapidly with increasing mass. The Salpeter Initial Mass Function is

${\displaystyle \xi (m)\Delta m=\xi _{0}\left({\frac {m}{M_{\odot }}}\right)^{-2.35}\left({\frac {\Delta m}{M_{\odot }}}\right).}$

where ${\displaystyle M_{\odot }}$ is the solar mass, and ${\displaystyle \xi _{0}}$ is a constant relating to the local stellar density.

Later authors extended the work below one solar mass (M_☉). Glenn E. Miller and John M. Scalo suggested that the IMF "flattened" (approached ${\displaystyle \alpha =0}$) below one solar mass.[7]

Pavel Kroupa kept ${\displaystyle \alpha =2.3}$ above half a solar mass, but introduced ${\displaystyle \alpha =1.3}$ between 0.08-0.5 M_☉ and ${\displaystyle \alpha =0.3}$ below 0.08 M_☉.

${\displaystyle \xi (m)=m^{-\alpha },}$

${\displaystyle \alpha =0.3}$ for ${\displaystyle m<0.08,}$

${\displaystyle \alpha =1.3}$ for ${\displaystyle 0.08<m<0.5,}$

${\displaystyle \alpha =2.3}$ for ${\displaystyle m>0.5}$

Chabrier 2003 for individual stars:

${\displaystyle \xi (m)=0.158(1/(\ln(10)m))\exp[-(\log(m)-\log(0.08))^{2}/(2\times 0.69^{2})]}$ for ${\displaystyle m<1,}$

${\displaystyle \xi (m)=km^{-\alpha }}$ for ${\displaystyle m>1,\alpha =2.3\pm 0.3}$

Chabrier 2003 for stellar systems (e.g. binaries):

${\displaystyle \xi (m)=0.086(1/(\ln(10)m))\exp[-(\log(m)-\log(0.22))^{2}/(2\times 0.57^{2})]}$ for ${\displaystyle m<1,}$

${\displaystyle \xi (m)=km^{-\alpha }}$ for ${\displaystyle m>1,\alpha =2.3\pm 0.3}$

There are large uncertainties concerning the substellar region. In particular, the classical assumption of a single IMF covering the whole substellar and stellar mass range is being questioned in favor of a two-component IMF to account for possible different formation modes of substellar objects. I.e. one IMF covering brown dwarfs and very-low-mass stars on the one hand, and another ranging from the higher-mass brown dwarfs to the most massive stars on the other. Note that this leads to an overlap region between about 0.05 and 0.2 M_☉ where both formation modes may account for bodies in this mass range.[9]

The possible variation of the IMF affects our interpretation of the galaxy signals and the estimation of cosmic star formation history[10] thus is important to consider.

In theory, the IMF should vary with different star-forming conditions. Higher ambient temperature increases the mass of clasping gas cloud (Jeans mass); lower gas metallicity reduces the radiation pressure thus make the accretion of the gas easier, both lead to more massive stars formed in a star cluster. The galaxy-wide IMF can be different from the star-cluster scale IMF and may systematically change with the galaxy star formation history.[11].

Measurements of the local Universe where single stars can be resolved is consistent with an invariant IMF[12] but the conclusion suffers from large measurement uncertainty due to the small number of massive stars and difficulties in distinguishes the binary systems from the single stars. Thus the IMF variation effect is not prominent enough to be observed in the local Universe.

Systems formed at much earlier times or further from the Galactic neighborhood, where star formation activity can be hundreds or even thousands time stronger than the current Milky Way, may give a better shot. It has been consistently reported both for star clusters[13] and galaxies[14] that there seems to be a systematic variation of the IMF. However, the measurements are less direct. For star clusters, the IMF changes over time due to complicated dynamical evolution. The study of galaxy-wide IMF is free from the dynamical effect since the spatially integrated light does not distinguish where a star is located. But it is very difficult to interpret those light signal. Different galaxy spectra may indicate an IMF difference but it could also be a difference in the age and metal abundance of the stellar population and a difference in dust extinction. To make it even more complicated, there are more than one and in fact, many stellar populations in a galaxy forming at different times. Thus it is almost always possible to construct a unique star formation history to explain the galactic spectrum without involving the IMF variation[15]. In this sense, the interpretation of the galactic observation cannot be exclusively IMF variation. However, the observations should be considered collectively rather than individually. When a large number of galaxy is observed and they show systematic spectral variation as a function of the current galaxy-wide star formation rate, it is more natural to interpret them as IMF variation due to the different of the star-forming environment instead of a fine set of star formation history of all those galaxies since star formation rate could fluctuate strongly in short timescales and has no correlation with the present-day star formation rate.