Skip to main content

Ram pressure stripping in high-density environments


Galaxies living in rich environments are suffering different perturbations able to drastically affect their evolution. Among these, ram pressure stripping, i.e. the pressure exerted by the hot and dense intracluster medium (ICM) on galaxies moving at high velocity within the cluster gravitational potential well, is a key process able to remove their interstellar medium (ISM) and quench their activity of star formation. This review is aimed at describing this physical mechanism in different environments, from rich clusters of galaxies to loose and compact groups. We summarise the effects of this perturbing process on the baryonic components of galaxies, from the different gas phases (cold atomic and molecular, ionised, hot) to magnetic fields and cosmic rays, and describe their induced effects on the different stellar populations, with a particular attention to its role in the quenching episode generally observed in high-density environments. We also discuss on the possible fate of the stripped material once removed from the perturbed galaxies and mixed with the ICM, and we try to estimate its contribution to the pollution of the surrounding environment. Finally, combining the results of local and high-redshift observations with the prediction of tuned models and simulations, we try to quantify the importance of this process on the evolution of galaxies of different mass, from dwarfs to giants, in various environments and at different epochs.

This is a preview of subscription content, access via your institution.

Fig. 1
Fig. 2
Fig. 3
Fig. 4
Fig. 5
Fig. 6
Fig. 7
Fig. 8
Fig. 9
Fig. 10
Fig. 11
Fig. 12
Fig. 13


  1. Ram pressure of the ICM comes from ions in the ICM (mainly protons), while the contribution from electrons is tiny (similar to the ICM viscosity). On the other hand, the thermal pressure of the ICM has roughly the equal contribution from free electrons and ions.

  2. The ram pressure Mira experiences only \(\sim 1.3 \times 10^{-11}\) dyn (assuming H i gas), which is only typical for the RP experienced by cluster galaxies at the cluster outskirts. Thus, an AGB star like Mira soaring in the ICM with a velocity of \(\sim \) 2000 km s\(^{-1}\) can have a tail of \(\sim \) 0.06 kpc.

  3. Different definitions of the virial radius were used in different works, e.g. \(r_{200}\), \(r_{180}\) and \(r_{101}\) (see, e.g. Bryan and Norman 1998 for detail).

  4. Note that many recent cluster works adopted the ICM density profile first proposed by Vikhlinin et al. (2006) (their equation 3). However, unlike the model proposed by Patej and Loeb (2015), this model is not physically motivated. Nevertheless, the best fit is also given here: \(E(z)^{-2} n_{\mathrm{e}} (x) = 0.00472 \Big (0.0818(\frac{x}{0.0735})^{-0.769} [1+(\frac{x}{0.0735})^{2}]^{-0.897} [1+(\frac{x}{0.280})^{3}]^{-0.570}\Big )^{1/2}\), good for \(x = 0.02 - 1.1\). Note that the central component of the Vikhlinin et al. (2006) model is not included as the density profile does not include the very central core.

  5. We did not use Eq. 13 in the following equations as we want to keep the mass dependency and Eq. 13 is not correct for real clusters.

  6. We recall that in such a configuration \(R_{\mathrm{0,star}} = 0.32\, R_{25}\), the effective radius (radius including half of the total galaxy luminosity), \(R_{\mathrm{eff}} = 0.59\, R_{25}\), and \(R_{\mathrm{0,gas}} = 0.61\, R_{25}\).

  7. As specified in Sect. 3.2.3, there are indications that clumpy structures such as giant molecular clouds decouple from the ram pressure wind.

  8. Kelvin–Helmholtz instabilities are due to the velocity difference across the interface between two fluids (e.g. Livio et al. 1980; Nulsen 1982; Mori and Burkert 2000). They manifest as small scales waves and vortexes as those observed in the atmosphere of Jupiter. Rayleigh–Taylor instabilities are due to the pressure of a light fluid pushing a heavy fluid, forming small-scale spikes and bubbles at their interface (e.g. Roediger and Brüggen 2008).

  9. For example, the radiative cooling time is 1.3–7.2 Gyr for several tail regions in NGC 4552’s X-ray tail beyond the galaxy, with the X-ray gas properties derived by Kraft et al. (2017).

  10. Different recipes for gas-to-dust ratio vs. metallicity relation are given in this reference.

  11. Diffuse filaments in the FUV and NUV bands must be interpreted with care since the UV emission could be originating from hydrogen 2-photon continuum or resonant lines of CIV and MgII generated by shocks rather than from stellar continuum emission (e.g. Bracco et al. 2020).

  12. The case of Mira indicates that the external pressure can perturb the wind produced by the mass loss of evolved stars as AGB stars (e.g. Li et al. 2019a). In general, stellar winds are subject to ram pressure.

  13. The H ideficiency parameter, first defined by Haynes and Giovanelli (1984), is a measure of the logarithmic difference between the expected and the observed H i mass of galaxies of different morphological type and size, where the estimated measure is taken from standard scaling relations of unperturbed galaxies in the field. For an updated calibration of these relations, see Boselli and Gavazzi (2009).

  14. Infall of the H i gas located in the outer regions to the inner disc has been invoked to explain the observed strong correlation between the total gas content (H i plus H\(_2\)) of isolated late-type galaxies, a large fraction of which is located outside the stellar disc, and their overall star formation activity (Boselli et al. 2001). The supply of this gas component to the inner star-forming regions, however, would occur on timescales relatively long compared to the typical ones for gas stripping (see Sect. 4.2).

  15. Note that the classical evaporation is assumed here. If instead saturated evaporation (Cowie and Songaila 1977) is assumed, the evaporation timescale typically increases by a factor of a few. On the other hand, galactic clouds are clumpy so the actual contact surface between the hot ICM and the cold clouds can be much larger than the surface of a single spherical cloud, which will reduce the evaporation timescale, along with the likely turbulence around the mixing layers.