Accelerating massive galaxy formation with primordial black hole seed nuclei

Abstract

If massive primordial black holes (PBHs) exist and constitute a fraction of the dark matter, they can dramatically catalyze galaxy formation. By acting as pre-existing, high-density seeds, they can shorten the galaxy assembly time to as little as 100 Myr for up to 10^8 solar mass PBH seeds, allowing for the rapid formation of host halos. Furthermore, low surface brightness or diffuse galaxies may represent a natural outcome of this process, perhaps as the residue of halos seeded by smaller PBHs that failed to accrete a major baryonic component.

Figures (6)

• Figure 1: The green line is the evolution in $r_e\sigma^2$ of 100,000 PBH particles that lose 1% of their mass per time step , to be compared with mass conserving particles (black line). Both axes are dimensionless. The x-axis is time ; the y-axis, $r_e \sigma^2$, velocity dispersion (v or $\sigma$) squared multiplied by the effective radius. The effective radius contains half the mass. The dimensionless mass loss rate dlogM/ dlogt is shown by the dotted blue line. It averages 19 (noted at the top right), which is higher than PBH emit for most of their lives. The red line shows that a nucleus with a mass as low as 300 particles energises the early time steps of a simulation. Uneven time steps were used , as shown by the distribution of the plus signs. The bumps in the curves are not numerical problems; they are an artefact of the sharp outer boundary of the initial uniform distribution of particles and correspond to the crossing time. The number of particles in the simulation is noted at the bottom right. The point to the right is a test of sensitivity to the initial distribution. Instead of uniform, a power law in radius was adopted. The point is placed at the end that run.
• Figure 2: The x-y, y-z, and x-z projection density contours of a 1% mass loss per timestep simulation. The contours are the loci of equal projected mass density, separated by 0.4 dex. Early formation creates halos that are more centrally concentrated. Elliptical contours are common. The color figure is velocity in the z direction from the x-y projection. Blue is approaching; the red points are redshifted. One unit in the color bar above is the v$_z$ velocity dispersion.
• Figure 3: Final radial profiles of simulations with supermassive PBH nuclei. These are calculated by azimuthally averaging the DM density. The colored numbers on the right are effective radii. The x-axis is pixels, the y-axis solar masses per pixel. Nuclei of 10$^5$, 10$^6$ and 10$^7$ M$_\odot$ are shown in black, red and green respectively. These are runs 53, 61 & 132 in Table A2. Although these simulations are scale free, once these masses are assigned to the PBH nuclei, pixel units and timescales follow. Effective radii are nominally in kpc. In blue are shown a 10$^6$ M$_\odot$ case with self-interacting DM. Adding a scattering cross section removes the cusp. This is not a property associated with PBHs.
• Figure 4: Final radial profiles of simulations with supermassive PBH nuclei as in the previous figure. A nucleus of 10$^5$ M$_\odot$ is shown in black (run 187). Effective radii, nominally in kpc, are also noted. In green and red we show 10$^4$ and 10$^6$ M$_\odot$ (runs 185 and 188) cases with an N $\sim$ M$^{-1}$ IMF. The colored numbers are effective radii. The difference between the cored profile and the other two is a gap in the mass function between the nucleus and the next PBH mass.
• Figure A1: The effect of mass around the nucleus on radial profile. These are all cusps with increasing density near the nucleus, with the exception of run 61, which has m$_1$ and m$_2$ five orders of magnitude below the nuclear mass.