The impact of magnetic fields during tidal disruption events

TL;DR

This work addresses how the magnetic field embedded in a star affects the early evolution of tidal disruption events. By combining global MHD simulations with a semi-analytic model, it maps when magnetic pressure becomes dynamically important and how it alters the debris stream, including rapid transverse widening and a magnetic–tidal equilibrium that can persist during the return to pericentre. A key finding is that if the initial stellar field exceeds roughly 10^4 G, magnetic pressure can magnetically support a substantial portion of the stream within about a year, significantly modifying the returning stream and potentially boosting radio emission from interactions with the ambient medium. The results provide physically motivated initial conditions for later phases of TDEs, with implications for disk formation, MRI development, jet production, and multiwavelength observables such as X-ray and radio signatures.

Abstract

During a tidal disruption event (TDE) the stream debris inherits the magnetic field of the star. As the stream stretches, the magnetic field evolves and can eventually become dynamically important. We study this effect by means of magnetohydrodynamic simulations and a semi-analytic model of the disruption of a main-sequence star by a supermassive black hole. For stellar magnetic fields stronger than $\sim 10^4\,\rm{G}$, magnetic pressure becomes important in a significant fraction of the mass of the stream, leading to a fast increase in its thickness, an effect that may impact its subsequent evolution. We find that this dynamical effect is associated with a phase of transverse equilibrium between magnetic and tidal forces, which causes the stream width to increase with distance to the black hole as $H \propto R^{5/4}$. In the unbound tail, this fast expansion could enhance the radio emission produced by the interaction with the ambient medium, while in the returning stream, it may qualitatively affect the subsequent gas evolution, particularly the gas dynamics and radiative properties of shocks occurring after the stream's return to pericentre. By characterizing the magnetohydrodynamic properties of the stream from disruption to the first return to pericentre, this work provides physically motivated initial conditions for future studies of the later phases of TDEs, accounting for magnetic fields. This will ultimately shed light on the role of magnetic fields in enabling angular momentum transport in the ensuing accretion disk, thereby affecting observable signatures such as X-ray radiation and relativistic outflows.

Figures (10)

• Figure 1: Fraction of the mass in the stream that becomes magnetically supported within time $t_{\rm target}$ after disruption, as a function of the initial magnetic field. The grey and black lines are obtained as described in Section \ref{['ss:Onset of magnetic pressure support']}. The pink lines correspond to the magnetic fields of main sequence stars inferred from asteroseismology measurements of red giants (see 2023AA...670L..16D), while the blue shaded region represents the magnetic field found in simulations of stellar mergers (Ryu_2025, Scheider2019).
• Figure 2: Snapshots showing the projection of the density and magnetic field on the orbital plane in the simulation initialized with a uniform field $B_{\star}=3\,\rm{MG}$ at times $t=0\, \rm{h}$, $t=1.5\, \rm{h}$ and $t=4\, \rm{h}$, corresponding respectively to distances from the black hole of $R=3\,R_{\rm t}$, $R=R_{\rm t}$ and $R=4.5\,R_{\rm t}$. The gray arrows point towards the black hole. As the star gets deformed and stretched the magnetic field aligns with the elongation direction. The colorbar describing the magnetic field strength is identical in the three panels.
• Figure 3: Left panel: evolution of accelerations $\ddot{H}_{\rm{m}}$, $\ddot{H}_{\rm{t}}$ and $\ddot{H}_{\rm{p}}$, for a parabolic stream element in the simulation initialized with a uniform magnetic field $B_{\star}=3\,\rm{MG}$. The purple dashed segments indicate the expected scalings for $\ddot{H}_{\rm{m}}$ and $\ddot{H}_{\rm{p}}$ during the early evolution, when the stream is still confined by self-gravity. The leftmost red vertical line marks the time $t_{\rm{mag}}$ at which $\ddot{H}_{\rm{m}}=\ddot{H}_{\rm{p}}$ in the simulation. After $t_{\rm{mag}}$ the stream element expands at constant speed causing $\ddot{H}_{\rm{m}}$ and $\ddot{H}_{\rm{t}}$ to follow the orange dashed segments. Later on, the tidal and magnetic accelerations reach an equilibrium with $\ddot{H}_{\rm{t}}=\ddot{H}_{\rm{m}}$ (rightmost vertical red line), after which they both follow the green dashed segments. Right panel: evolution of $\ddot{H}_{\rm{m}}$ and $\ddot{H}_{\rm{t}}$ for stream elements of boundness $\mu=-1.1$ and $\mu=1.1$ in the simulation initialized with a uniform magnetic field $B_{\star}=10^5\,\rm{G}$. The vertical red solid line marks the transition time $t_{\rm{tr}}$ (equation (\ref{['eq:t_tr']})), while the vertical red dashed line indicates the time to reach apocenter $t_{\rm{apo}}$.
• Figure 4: Projection of the density on the orbital plane for the hydrodynamic (top) and magnetized (bottom) tidal stream with initial field $B_{\star}=3\,\rm MG$ at time $t\approx 442\, \rm{h}$. The grey arrow points towards the black hole. The inset plot shows a zoom-in of the near-parabolic stream section ($|\mu|<0.05$), which at this point of the simulation is located at $163 \, R_{\rm{t}}$ from the black hole. The width of the magnetized stream here is a factor of a few larger than in the hydrodynamic case.
• Figure 5: Evolution of the stream vertical width as a function of the distance from the black hole in the simulation for different orbital energies corresponding to $\mu=-1.3$ (blue lines), $\mu=0$ (purple lines), and $\mu=1.3$ (grey lines). The dashed lines correspond to a simulation without magnetic fields, while the solid lines correspond to magnetized cases with initial magnetic field $B_{\star}=3\,\rm MG$. The black dashed segments indicate the expected scalings for the different parts of the evolution.