Magnetic Flux Emergence in Binary Neutron Star Remnants

TL;DR

This study addresses whether magnetic flux can emerge from a post-merger binary neutron star remnant to power magnetically driven jets. It employs high-resolution GRMHD simulations with a twisted toroidal flux tube embedded in a TOV remnant modeled with the DD2 EOS and an approximate neutrino-trapping treatment to quantify emergence conditions and timescales. The key finding is that emergence occurs only for extremely strong fields near or above $10^{17}$ G; more typical post-merger fields around $10^{16}$ G are dominated by hydrodynamic buoyancy and do not emerge effectively, implying that remnant-driven jets are unlikely and that jets may originate in the disk instead. These results highlight the crucial roles of toroidal field tension and EOS in magnetic buoyancy and motivate future global, high-resolution simulations exploring different EOSs, rotation, and field geometries to fully assess jet-launching mechanisms in BNS mergers.

Abstract

Using high-resolution AthenaK simulations of a twisted toroidal flux tube, we study the flux emergence of magnetic structures in the shear layer of a hot massive neutron star typical of a binary neutron star remnant. High-resolution simulations demonstrate that magnetic buoyant instabilities allow for emergence only for extremely large magnetic fields significantly exceeding $10^{17}~\mathrm{G}$, and more typical fields around $10^{16}~\mathrm{G}$ are instead dominated by hydrodynamic effects. Because merger remnants tend to be stable against hydrodynamic convection, our work places strong limitations on the mechanisms by which massive binary neutron star remnants can produce the magnetically-driven outflows needed to power jets.

Figures (12)

• Figure 1: A slice of the $r$-$\phi$ plane at $t=0.39~\textrm{ms}$ comparing $|B|$ from H1e16_5 (left) to the tracer scalar $X$ from H0_5 (right). The dashed black line marks the surface of the star, and the dot-dashed green lines mark the boundaries of the initial position of the flux tube. The nearly identical evolution indicates that the dynamics of H1e16_5 are driven primarily by hydrodynamic buoyancy from the original perturbation rather than magnetic buoyancy.
• Figure 2: The magnetic energy $E_B$ (solid lines) and kinetic energy $E_K$ (dashed lines) as a function of time normalized by the total energy for all the runs with perturbations. The dotted vertical line marks the emergence time for H7e17_5.
• Figure 3: Snapshots of $\tilde{B}^\phi$ in the $r$-$\phi$ plane from H1e17_5 at $t=0.39~\textrm{ms}$ (left) and $t=1.31~\textrm{ms}$ (right). The location of the flux tube is roughly the same in both, but some expansion and diffusion is apparent at the later time. The dashed black line indicates the location of the stellar surface from the TOV solution, and the dot-dashed green lines mark the boundaries of the flux tube at $t=0$.
• Figure 4: The magnitude of the magnetic field $|B|$ in the $r$-$\theta$ plane from H7e17_1_5 at $t=0~\textrm{ms}$ (top-left), $t=0.15~\textrm{ms}$ (top-right), $t=0.25~\textrm{ms}$ (bottom-left) and $t=0.34~\textrm{ms}$ (bottom-right). The contours in black, blue, and red denote $10^{15}~\textrm{G}$, $10^{16}~\textrm{G}$, and $10^{17}~\textrm{G}$, respectively. As the flux tube rises, it expands into the umbrella shape characteristic of a stratified environment with a decreasing scale height. The dashed black line marks the location of the stellar surface from the TOV solution.
• Figure 5: Streamlines of the magnetic field for H7e17_5 (top) and H1e17_5 (bottom) at the initial time (left) and a late time (right). The surface contour marks $\rho=10^{-5}~\textrm{M}_\odot^{-2}\approx6.2\times10^{12} \mathrm{g}/\mathrm{cm}^3$, which is the approximate location of the surface. Note that the color scale is normalized to the maximum field strength in each run. H7e17_5 expands rapidly as it emerges, while H1e17_5 remains trapped beneath the surface and largely intact.