Bases del Modelo Integral de Reinserción Social

1.2.1

Low mass star formation

Although it is not within the remit of this thesis, it is helpful to give a brief overview of the generally accepted formation mechanism for low-mass stars, es- pecially since some theories paint massive star formation as a scaled-up version of this mechanism (see subsection 1.2.3).

A seminal review by Shu et al. (1987) tied together and summarised previous works on low-mass star formation, and concluded that this process starts with slowly rotating density enhancements (‘cores’) within molecular clouds. These cores are either sub-critical, whereby the internal pressure support the cores against collapse, or super-critical, whereby the cores cannot support themselves and collapse due to self-gravity. In the sub-critical case, turbulent and magnetic (through ambipolar diffusion) support slowly leak out of the core until it turns super-critical. It is important to note that either way, collapse proceeds from the inside-out. Due to the overall angular momentum vector of the initial rotating core, the collapse forms an accretion disc through which accretion proceeds upon

a central protostar. Stellar winds/jets drive molecular outflows out of cavities formed along the poles of the rotation axis, which eventually dissipate the sur- rounding circumstellar envelope as they widen, leaving a remnant disc from which planets may eventually form.

Due to the relative flux contributions at various wavelengths from separate components of different temperatures, the spectral energy distribution can be used to establish an evolutionary classification system. Measurement of the spec- tral index between 1 µm and 20 µm allows low-mass protostars to be categorised as class 0 (undetectable at λ > 10 µm; Andre et al. 1993), class I, class II or class III _{(0 < α < 3, −2 < α < 0 or −3 < α < −2 respectively; Lada 1987). The} reader is directed to Figure 1 of Feigelson & Montmerle (1999) and Figure 1.1 of Isella (2006) for a complete description and illustration of each evolutionary class.

1.2.2

Massive star formation

The astrophysical community remains divided between the prominent paradigms of massive star formation. On one hand we have the idea that it is the core that collapses to form a massive star of a defined fraction (‘star formation efficiency’) of its mass, and on the other is the idea that the massive star’s mass can exceed that of its core by sourcing material from the surrounding clump. More extreme processes involving the collision of two forming stars to form a protostar of higher mass have also been considered.

1.2.3

Turbulent core model

Essentially the turbulent core model (McKee & Tan 2002, 2003) is similar to the accepted mechanism for low-mass star formation (see subsection 1.2.1), whereby the material a massive star accretes during its formation comes directly from the initial prestellar core it evolves from. Observationally-speaking this translates into a similarity of the core mass function, or CMF, (distribution of prestellar core masses as a function of a mass) with the IMF (as well as a stellar mass to core mass ratio of ∼ 0.5, Tan et al. 2014). This approach conserves angular momentum in the form of a rotating, flattened, accretion disc and associated bi-polar jets. Direct observations of ‘massive rotating structures’ (i.e. potential accretion discs) of comparable mass to the accreting star (upto ∼ 20 M⊙) are

commonplace (see Cesaroni et al. 2007, for a review) supporting this view. Disc- jet systems have been observed in a handful of examples in the mm/cm regime, (see subsection 1.4.4) however when it comes to the most massive stars (O4−O8 type), only toroids of material around the central star have been observed (e.g. NIR observations of W33a by Davies et al. 2010). This may be due to the small sample size of O4−O8 type stars, their distance, the relative luminosity of the tori to discs, or other effects. Criticisms on the basis of the increased radiative forces on the disc of gas and dust from the high radiative output of an MYSO do produce a problem of accretion being counteracted by radiation pressure (Zinnecker & Yorke 2007). It is thought this issue can be alleviated as the disc forms only a small cross-sectional area (i.e. non-spherical accretion/non-Bondi Hoyle), to the radiation’s direction (Krumholz et al. 2009). This means only exposing a small fraction of the disc to its full force, and thus avoiding dispersion.

1.2.4

Competitive Accretion

Simulations by Bonnell et al. (1997), who considered a strongly gravitationally- bound clump, showed that the protostars in the central regions had much higher accretion rates than average, spawning the idea of competitive accretion. This process operates on two main principles, the first being that protostars at the centre of clumps will become the largest, as there is simply more matter (denser environments) within their accretion domains. The second principle is that those protostars that have formed first, have more time to accrete and hence are grav- itationally more attractive (i.e. larger accretion radius) compared to their neigh- bourhood siblings.

In comparison to core accretion, the material that forms the protostar is siphoned from all parts of its natal clump (‘clump-fed’) as the accreting MYSO moves internally through the clump. The knock-on effect is that the core mass function should no longer mirror the initial mass function (Tan et al. 2014). Other predictions include changing accretion disc rotation axes (due to chaotic mass flows to the forming star and dynamical interactions in crowded environments), which observationally should present itself in the precession of outflow directions. Massive stars should also be seen to always form at the centre of clumps and with relatively small accretion disks.

However, due to the Bondi-Hoyle process which this model employs, radia- tive forces impeding accretion become a problem at M⋆ > 10 M⊙ (Tan et al.

2014). Outflows, regularly observed towards sites of massive star formation (see section 1.3), further exacerbate this issue by halting accretion along their axis of motion. Indeed, simulations of this mechanism have yielded poor accretion

rates of 4.6 × 10−5_M

⊙yr−1 for even the most massive (46 M⊙) star formed in

the simulated clump (Wang et al. 2010), which are an order of magnitude too small. Possible solutions include a higher accretion efficiency, in order to reduce the time for massive stars to form, or time-variable accretion rates (such as the accretion-driven bursts seen in the simulations of Meyer et al. 2017) since those shown in Wang et al. (2010) were relatively constant. This last possibility has intriguing implications on the mass loss of collimated jets, since accretion and mass loss mechanisms are generally linked (see subsection 1.4.1).

1.2.5

Stellar Mergers

Also referred to as coalescence in the literature, this process essentially forms massive stars by the physical merging of multiple, lower-mass precursors (Bon- nell et al. 1998). Originally it was a concept designed to overcome the radiation pressure problem, however it was found to be an unlikely candidate for the dom- inant process of MSF due to the high stellar densities required for significant merger rates (∼ 108_pc−3_{). As such it shall not be discussed in any more detail,}

however it has been included here for completeness and may still contribute to the formation of the highest mass stars.