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tokens: of the disc will be In our model, the relative change in disc mass and accretion rate in an encounter is determined by γ and the truncation radius relative to the disc characteristic radius just before the encounter, . If the disc is truncated, its viscous time-scale decreases and the disc starts evolving faster. We ignore the orientation of the discs, so the equation for truncation radius is the average truncation radius over all disc inclinations. 2.4 Numerical implementation We use the fourth-order Hermite N -body code ph4 , which is incorporated into the AMUSE framework. The softening length is and the time-step parameter is η = 10 −2 . The simulations start in virial equilibrium and last for 2 Myr. Collisional radii for each star in the cluster are defined depending on the stellar mass and the theoretical size of the disc evolving in isolation, given by equation ( 1 ). The initial collisional radius for each star is given by the distance at which the most massive star in the cluster can truncate its disc. The collisional radii of all stars are updated every 2000 yr to account for the viscous evolution of the discs. Two stars are considered to be in an encounter if the distance between them is less than the sum of their collisional radii. The encounter distance between the stars is then computed by analytically solving the two-body problem. If the encounter is strong enough as to truncate the discs, disc parameters of both stars are updated as described in Section 2.3 . After each encounter, the collisional radii of both stars are set to 0.49 r enc in order to prevent the N -body code from detecting the same encounter multiple times during the same 2000 yr window. 3 RESULTS 3.1 Initial conditions 3.1.1 Cluster properties Simulations were run for clusters with 1500 stars. The stellar masses are randomly sampled from a Kroupa initial mass distribution (Kroupa 2001 ) with lower mass limit 0.1M ⊙ and upper mass limit 100M ⊙ . The stars are initially distributed in a Plummer sphere (Plummer 1911 ) with Plummer radius 0.5 parsec. All stars are assumed to be single and coeval. All simulations were carried out for . We define three different gas scenarios for the simulations: No gas Gas presence Gas expulsion starting at Each simulation was run for two values of the turbulence parameter: α = 10 −2 and α = 5 × 10 −3 . 3.1.2 Disc properties Based on the observations of low-mass stars carried out by Isella et al. ( 2009 ) and the estimations obtained by Hartmann et al. ( 1998 ), the initial characteristic radius R c of the circumstellar discs was chosen to be with . Circumstellar discs with masses larger than 10 per cent of the mass of their star are the most likely to be gravitationally unstable (Kratter Lodato 2016 ). According to hydrodynamical simulations of collapsing gas clouds, the disc-to-star mass ratios of embedded protostars are large enough for gravitational instabilities to occur (Vorobyov 2011 ), but these instabilities lead to