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I'm reading material that is seemingly contradictory. Some sources indicate that the evolution of a protostar to a main sequence star is characterised by a stellar wind that precludes the accretion of further in-falling material. That is, the (young) star now has a constant mass. However, other sources suggest that material may continue to accrete for a (brief) period after the protostar has become a main sequence star.
Can someone please confirm the actual process?
At the end of the protostar phase a vigorous outflow from the star develops called the T-Tauri Wind and this could cut off accretion. Eventually, it develops into a normal star and the strong wind dies down. Material that was not totally blown away could then continue to fall down and be accreted
Accretion of in-falling material for a young main sequence star - Astronomy
We present the initial result of a large spectroscopic survey aimed at measuring the timescale of mass accretion in young, pre-main-sequence stars in the spectral type range K0-M5. Using multi-object spectroscopy with VIMOS at the VLT we identified the fraction of accreting stars in a number of young stellar clusters and associations of the ages of between 1-30 Myr. The fraction of accreting stars decreases from
2% at 10 Myr. No accreting stars are found after 10 Myr at a sensitivity limit of 10 -11 M ⊙ yr -1 . We compared the fraction of stars showing ongoing accretion (f_acc) to the fraction of stars with near-to-mid infrared excess (f_IRAC). In most cases we find f_acc <f_IRAC, i.e. , mass accretion appears to cease (or drop below detectable level) earlier than the dust is dissipated in the inner disk. At 5 Myr, 95% of the stellar population has stopped accreting material at a rate of ⪆10 -11 M ⊙ yr -1 , while
20% of the stars show near-infrared excess emission. Assuming an exponential decay, we measure a mass accretion timescale (τ_acc) of 2.3 Myr, compared to a near-to-mid infrared excess timescale (τ_IRAC) of 3 Myr. Planet formation and/or migration, in the inner disk might be a viable mechanism to halt further accretion onto the central star on such a short timescale.
Based on observations collected at the European Southern Observatory, Paranal, Chile (Proposal ID: 078.C-0282 081.C-0208).
Accretion of in-falling material for a young main sequence star - Astronomy
Context: The brightness of FUors increases by several magnitudes within one to several years. The currently favoured explanation for this brightness boost is that of dramatically rising accretion from the disc material around a young star. The mechanism leading to this accretion increase is a point of debate.
Aims: Choosing the Orion nebula cluster as representative, we simulate accretion bursts driven by encounters in dense stellar environments. We investigate whether properties like rise and decay times, event frequency, etc., speak for encounters as a possible cause for FUor phenomena.
Methods: We combine cluster simulations performed with the Nbody6++ code with particle simulations that describe the effect of a fly-by on the disc around a young star to determine the induced mass accretion.
Results: The induced accretion rates, the overall temporal accretion profile, the decay time, and possibly the binarity rate we obtain for encounter-induced accretion agree very well with observations of FUors. However, the rise time of one year observed in some FUors is difficult to achieve in our simulations unless the matter is stored somewhere close to the star and then released after a certain mass limit is transgressed. The severest argument against the FUors phenomenon being caused by encounters is that most FUors are found in environments of low stellar density. We extend the discussion to eccentric binaries and gravitationally unstable discs and find that both models have similar problems in achieving the necessary rise times.
Conclusions: We find no conclusive answer as to whether the observed FUors are triggered by encounters. However, it seems an intense accretion burst phase should exist - possibly an FU phase - early on in the development of dense clusters. We predict that in dense young clusters these outbursts should happen predominantly close to the cluster centre and with high mass ratios between the involved stars.
Accretion onto Pre-Main-Sequence Stars
Accretion through circumstellar disks plays an important role in star formation and in establishing the properties of the regions in which planets form and migrate. The mechanisms by which protostellar and protoplanetary disks accrete onto low-mass stars are not clear angular momentum transport by magnetic fields is thought to be involved, but the low-ionization conditions in major regions of protoplanetary disks lead to a variety of complex nonideal magnetohydrodynamic effects whose implications are not fully understood. Accretion in pre-main-sequence stars of masses ≲1M⊙ (and in at least some 2–3-M⊙ systems) is generally funneled by the stellar magnetic field, which disrupts the disk at scales typically of order a few stellar radii. Matter moving at near free-fall velocities shocks at the stellar surface the resulting accretion luminosities from the dissipation of kinetic energy indicate that mass addition during the T Tauri phase over the typical disk lifetime ∼3 Myr is modest in terms of stellar evolution, but is comparable to total disk reservoirs as estimated from millimeter-wave dust emission (∼10 −2 M⊙). Pre-main-sequence accretion is not steady, encompassing timescales ranging from approximately hours to a century, with longer-timescale variations tending to be the largest. Accretion during the protostellar phase—while the protostellar envelope is still falling onto the disk—is much less well understood, mostly because the properties of the central obscured protostar are difficult to estimate. Kinematic measurements of protostellar masses with new interfometric facilities should improve estimates of accretion rates during the earliest phases of star formation.
Accretion of in-falling material for a young main sequence star - Astronomy
Figure 1: Two images of V1647 Orionis and McNeil&rsquos Nebula. The image on the left is an optical color composite taken about four years ago with GMOS-North on UT 2004 February 14. The image on the right is also an optical color image taken about one year ago on UT 2007 February 22.
Figure 2: Expanded view of the 2.12-2.35 micron region of the near infrared spectroscopy of V1647.
Figure 3: Plot of the 8-13 micron silicate absorption band optical depth extracted from the mid infrared spectrum.
Figure 4: Optical spectroscopy of V1647 Orionis from GMOS-North obtained on UT 2007 February 22.
A &ldquonew&rdquo star appeared in the constellation of Orion in late 2003 when the young pre-main sequence star V1647 Orionis went into outburst. The eruption and huge increase in brightness of the object resulted in the appearance of a reflection nebula called &ldquoMcNeil&rsquos Nebula,&rdquo named after the amateur astronomer, Jay McNeil, who discovered the object and alerted the world.
During the outburst the star and nebula remained bright for approximately 18 months before fading rapidly over a six month period. By early 2006 the star and its environment were very similar to their pre-burst stage. The event was monitored and observed with many ground- and space-based facilities and Gemini Observatory played a key role in monitoring the event during its eruptive and quiescent phases. A team led by Colin Aspin (IfA/University of Hawai&lsquoi), Tracy Beck (STScI) and Bo Reipurth (IfA/University of Hawai&lsquoi) spearheaded the monitoring campaign of this unique event.
The eruption of V1647 Orionis is most likely associated with a mass dumping of the inner regions of a heated circumstellar disk onto the young stellar photosphere. The spectacular flaring in brightness of the object is due to a significant increase in accretion luminosity and the clearing or destroying of surrounding dust by an energetic wind that made the star visible. These eruptions are thought to be repetitive and indicative of periods when a significant fraction of the final star&rsquos mass is accreted.
The authors describe three phases for the V1647 Orionis latest eruption:
- Before November 2004 is the pre-outburst phase
- From November 2004 to February 2006 is the outburst phase
- From February 2006 is the quiescent phase
The Gemini observing campaign led by Aspin has revealed some interesting results, particularly for the quiescent period. These include:
- McNeil&rsquos Nebula is faintly visible in these GMOS-N images (Figure 1 right) indicating that the nebular material is still weakly illuminated by the star V1647 Orionis. At the time of acquisition of the GMOS-N imaging and spectroscopic data , V1647 Orionis had an r&rsquo magnitude of 23.3.
- NIRI spectroscopy has revealed for the first time in this type of object the presence of molecular overtone absorption from CO and other key diagnostic atoms like Na and Ca (possibly betraying the photosphere of the star), see Figure 2. The 2um spectroscopy shown in the paper is from IRTF not NIRI. We did publish NIRI spectroscopy but from just after the outburst, not in quiescence.
- The star has a mass of about 0.8 solar mass and its age is about half a million years or less.
- V1647 Orionis in this pre-main sequence phase is about five times more luminous than the Sun.
- Material is falling onto the star at a rate of about one millionth of a solar mass per year.
- Mid infrared observation with MICHELLE/Gemini show evidence of silicate dust evolution over the outburst-to-quiescence period, see Figure 3.
In a previous article on V1647 Orionis, Aspin studied a previous outburst of the star which occurred in 1966. It seems that perhaps V1647 Orionis &lsquowakes up&rsquo every 37 years but soon (after 1 to 2 years) tires and takes another long nap!
For more details, read the article "V1647 Orionis: One year into quiescence", by C. Aspin, T. Beck and B. Reipurth in The Astronomical Journal, January 2008, pp. 423-440.
For more details on the 1966 outburst of V1647 Orionis, read the article "The 1966-1967 Outburst of V1647 Orionis and the Appearance of McNeil's Nebula", by C. Aspin and others in The Astronomical Journal, Volume 132, Issue 3, pp. 1298-1306.
Accretion of in-falling material for a young main sequence star - Astronomy
Young Stellar Objects (YSOs) are stars in the first phase of theirs lives, before they enter the main sequence of the Hertzsprung-Russell diagram and are fed by stably hydrogen fusion. YSOs are formed by contraction (and fragmentation) of molecular clouds. Contraction can be started by a variety of factors, such as general density fluctuations in the interstellar medium, radiation pressure of nearby stars, or shock waves of supernova events that lead to local compressions. The contraction of the molecular cloud is driven by gravity, the cloud actually collapses in free falling. The gravitational energy is released by radiation and in turn influences the collapse by its radiation pressure, which counteracts gravity. The dense center of the molecular cloud is the new protostar. A protostar emits light due to the heat released by the gravitational collapse. Its core temperature, however, is still too low to maintain nuclear fusion. In this protostar stage, the star is still growing by mass accretion from the surrounding molecular cloud, which lasts until either the entire cloud is incorporated or until the radiation pressure of the new star is powerful enough to blow off the remainders of the cloud.
Protostars, proto-stellar disks, jets and Herbig-Haro-Objects
Due to the preservation of angular momentum, the molecular cloud cannot simply collapse. Instead, a protostellar disk is formed around the protostar during contraction of the cloud. Along the axis of rotation, in-falling material has only little angular momentum and the in-fall proceeds relatively unhindered. Therefore, the molecular cloud becomes thinner along the axis of rotation and two cone-shaped voids are formed at the poles, which allows light from the star to escape and to illuminate these cones from the inside. Depending from the viewing angle, we see the molecular cloud illuminated by the young star as a bipolar nebular (viewed from the side), as a fan-shaped nebula (at a low angle from above the disk) or in a crescent or even ring shape with increasing viewing angle.
Material migrates within the protostellar disk towards the star by internal friction: the protostar accretes material.
This image shows an artistic view of the dusty protoplanetary disk around a massive young star.
The spinning-up of protostar and protostellar disk winds up magnetic vortices, leading to strong magnetic fields along the polar axis and the formation of bipolar outflows or jets. These jets may hit the surrounding interstellar medium or the remainders of the collapsing molecular cloud, leading to strong shock fronts, so-called Herbig-Haro objects (HHs).
Relatively bright Herbig-Haro objects (HH1 and HH2) can be found, for instance, in a molecular cloud south of the Key Hole Nebula NGC 1999 in Orion at the lower edge of this image:
Patrick Hartigan at Rice University in Houston succeeded in detecting the dynamics within the jets and shock fronts of several HHs using the Hubble Space Telescope. Movies of these dynamics are available on his web pages.
Bewegung des Jets von HH 1, Patrick Hartigan
Further on on the evolution/contraction of the protostar, the mechanism of transport of the released gravitational energy out of the core of the protostar switches from convection to radiation. This leads to more efficient cooling of the core, which is important in particular for the heavier protostars, as it allows contraction to move on more rapidly. During this process, the star is found in the Hertzsprung-Russell diagram above the main sequence, moving downward along the so-called Hayashi line. Finally the core of the new star becomes hot and dense enough to maintain stable hydrogen fusion: A new star is born.
Herbig Ae/Be, T Tauri and FU Ori stars
During this transition stage, the new stars have not yet reached a stable hydrostatic equilibrium. Instead, they further contract despite that hydrogen fusion may have started. The new stars have not yet reached the main sequence, but are placed still above it. They are still larger and therefore brighter than main sequence stars of same temperature (and hence same spectral class). During this transition stage, the new stars are quite variable and may also show strong bursts of brightness (flares). These bursts reflect the erratic infall of material from the accretion disk onto the star. By excitation of the thin outer atmosphere, the stars display emission lines during this phase of their lives.
Stars smaller than 2 solar masses are classified as T Tauri stars, after their prototype, while the heavier ones are called Herbig Ae/Be stars (e stands for emission lines). The stars are identified by several criteria: presence of emission lines (in particular the Balmer series of hydrogen), an infrared excess of their radiation due to the enveloping dust in the surrounding disk, and their situation in a star forming region. The latter is verified by the projected location (for instance in a dark cloud) and the presence of a reflection nebula associated with the star, which secures the location of the star within the molecular cloud. The sub-types of YSOs are distinguished by spectral classification and their mass (B and A for Herbig Ae/Be stars and F, G, K, or M for T Tauri types). The pre-main sequence phase is in both cases short as compared with the entire life span of the star and lasts 1 to 10 million years for the massive Herbig Ae/Be stars to 10 to 100 million years for the less massive T Tauri stars.
T Tauri stars have a further sub-group termed FU Ori stars (or short "Fuors"), which due to erratic accretion are subjected to flares and very pronounced bursts of brightness of up to 6 magnitudes. It is conceivable that FU Ori-like behavior represents a phase during the development of most T Tauri stars. The FU Ori type is hence not a type of stars of its own, but rather represents a temporary stage during the evolution of a T Tauri star. In addition, there is also the Exor type (after EX Lupi), showing flares at shorter time scales.
Evolution of Young Stellar Objects
The evolution of a YSO is classified into four subsequent phases, which correlate more or less with the increasing exposure of the star in its envelope and the resulting changes of its spectrum (this classification scheme is actually based on the spectral energy distribution (SED) of the YSO).
During the first phase, the molecular cloud collapses to a protostar, that remains completely buried in its surrounding envelope, eluding direct observation in visible light. Only the infrared radiation of the warm envelope can be observed. This stadium of gravitational collapse is going along with a tremendous increase of the rotational frequency of the star and the accretion disk (spin up). In the following, the star becomes exposed and the spectrum of the YSO shows the black body spectrum of the emerging star, superimposed with the still substantial infrared excess of the envelope at lower frequencies. These stages are the classical T Tauri star stage (CTTS), with an active disk, accompanied with formation of jets and Herbig-Haro objects, and the weak-line T Tauri star stage (WTTS), with a passive disk, the onset of nuclear fusion and fragmentation of the protoplanetary disk. During the final stage, the star moves onto the main sequence with continuous hydrogen burning, the protoplanetary disk fragments into a planetary system (carrying most of the angular momentum of the YSO), and the infrared excess in the spectrum vanishes.
Image of the dust disk around the young star IM Lupi (upper) and other YSOs (lower). Images are from SPHERE, ESO's Exoplanet Research Instrument at VLT in Chile.
Pre-main sequence stars are hence surrounded by (dusty) accretion disks and the (dusty) remainders of the molecular cloud in which they were formed. The accretion (or protoplanetary) disks are the cradles of new planets around the young star. The disk may often completely hide the star or at least strongly attenuate its light, thereby preventing its direct observation. Along the polar axis with its cone-shaped voids, however, the light of the star can pass relatively unhindered. In the telescope, this light can often be seen as it is illuminating the surrounding molecular cloud where it is scattered by dust particles. The bipolar appearance of the reflection nebula is hence a direct consequence of the presence of an accretion disk, which cuts the escaping star light down to narrow cones of light.
Variable nebula and polarization
Many of these reflection nebula are highly variable, which may have several reason:
A very impressive example of such casting of shadows is displayed by Hubble's Variable Nebula, NGC 2261, in the constellation of Monoceros. The animated image below shows a unique series of images by Tom Polakis from Tempe/Arizona, who photographed NGC 2261 over several weeks. The animation clearly shows the propagation of light and shadows along the reflection nebula.
variability of Hubble's Variable Nebula NGC 2261, Tom Polakis
Other reflection nebula with highly variable brightness are Gyulbudaghian's Nebula around PV Cephei , McNeil's Nebula in the molecular cloud of M78 in Orion, and the Nebula around Z Canis Majoris . The reflection nebula of PV Cephei, that used to be an easy target a few years ago with medium-sized telescopes, had been very difficult even with my 22" Dob over several years, until it brightened again in 08/2013.
Variability of the RN of PV Cephei on DSS red plates (compare Adam Block's image at the very of 2008)
McNeil's Nebula around V1647 Ori was discovered in 2004 by amateur Jay McNeil photographically. By 2006 it had dimmed again beyond visual detection. In 2008, it brightened again and had stayed on a relatively high level since then (this being the state of things until 2011), being accessible for telescopes around 18 inches or up.
McNeil's Nebula in 2006 (lower panel) and 2011 (upper panel).
All three stars, Z CMa, V1647 Orionis, and PV Cephei, are FU Ori variables with flares (with PV Cephei with its flares on shorter time scales, last flares were at 2004 and 2013, being rather an EXor). Another nebula that had been known for its brightness fluctuations on historical time scales is Hind's Variable Nebula around T Tauri.
A further phenomenon is the polarization of scattered light. Polarization of light can be verified by the appropriate filters, similar as described for the protoplanetary nebula . These have, despite the name, nothing in common with protoplanetary disks, but are formed during the late phases of stellar evolution prior to formation of a full planetary nebula. Similar polarization-dependent observations of YSOs are naturally restricted to the brightest objects.
Formation of planets around YSOs
In a 2014 press release , ALMA has resolved the protoplanetary disk around the T Tauri star HL Tauri at 1.2mm wavelength into dense rings and gaps. These gaps presumably report accretion of disk material by evolving planets within the disk.
HL Tauri is situated in the same molecular cloud as Sharpless 239, and in its direct vicinity several other YSOs can be found:
Visual observation of YSOs is like entering new territory. They are just too exotic and furthermore mostly very dim. There are only few well-known and brighter objects. The most well-known representatives are certainly Hubble's Variable Nebula and NGC 1999, which are rewarding objects, displaying structure already in smaller amateur telescopes. Further objects within reach of medium-sized telescopes are Ced 62 (NGC 2163) and Parsamian 21. Most other objects require large telescopes and are even then difficult targets. In many cases, only a stellar object can be discerned, the pre-main sequence star itself. The surrounding reflection nebula, so they are observable, are often extremely faint or outshined by the star. Their sometimes bizarre structure, which is very prominent on the DSS images, can be seen visually only in few cases. While it is usually not a problem to distinguish between star and surrounding nebulosity, the further distinction between reflection nebula, jet, and Herbig-Haro object is often difficult or not unequivocal. In particular the HHs are (despite few exemptions) very small and extremely faint. Nevertheless, these are very interesting objects, and it is exciting to observe stars in these very early stages of their lives (which is also the stage of the development of planets). And due to their high intrinsic variability, you never know what to expect!
The Young Stellar Objects Observing Guide
This observing guide introduces more than fifty pre-main sequence stars with surrounding reflection nebula with DSS images, finder charts, and observing reports at the eyepiece of my 22" Dob.
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A dense shell forms around the shocked core structure
As shown in Fig. 1 (B to D), what is observed in the laboratory and substantiated by the two MHD (in 2D and 3D) simulations (see Materials and Methods) is consistent with the following scenario: On impact, the stream, halted by the obstacle, induces the formation of an inward shock and of a reverse shock within the stream itself. The front of the reverse shock is localized by the density jump observed at the edge of the central core in Fig. 1 (B to D). As shown in the Supplementary Materials, we verify that, in the laboratory, the increase of the plasma electron density in the core postshock region corresponds closely to what is expected from the Rankine-Hugoniot strong-shock conditions (28).
Simultaneously, driven by the thermal pressure, which locally overcomes the magnetic pressure (β > 1), the highly conductive and shock-heated core plasma expands radially outward, compressing and distorting the magnetic field lines (Fig. 1D, white lines). Lateral expansion is then halted by the increasing magnetic field, which redirects the plasma flow toward the stream and forms an enveloping structure that we denote here as the “shell.” In Fig. 1 (B to D), we observe the ejected flow, the shell quickly overtaking (along z) the propagating reverse shock due to the longitudinal redirection of the ejected flow.
We stress that the plasma β of the core overwhelmingly determines the shell formation. For the regimes investigated here, radiative losses and thermal conduction play an obvious role in determining the details of the thermodynamic properties of the laboratory and astrophysical plasmas and, for example, their spectral signatures. However, as long as the core plasma β are similar (30 and 107 for the laboratory and astrophysical plasmas, respectively see table S2), the dynamics of the formation shell is largely insensitive to these details.
We also note that the laboratory accretion streams are sustained and interact with the obstacle over a time scale that is long compared to that of the formation of the shell (that is, those shown in Fig. 1C). Nevertheless, the absence of a gravitational field capable to exert an influence on the plasma hampers our ability to experimentally reproduce the long-term dynamics of an idealized astrophysical accretion column (25). Instead, the experiments are limited to the initial impact and formation of the plasma shell when gravity does not have an appreciable role (which is valid for the first two frames of Fig. 1, C and D see Materials and Methods). This is confirmed by comparing the simulations without gravity presented in our paper (Fig. 1D and movie S2) with long-term astrophysical MHD simulations that include gravity (25). Such comparison shows that the initial dynamics and shell formation are qualitatively similar and are largely unaffected by gravitational forces.
Our simulations do not account for radiative transfer effects, as detailed in the “Synthesis of the x-ray emission and comparison with astrophysical objects” section. This assumption can be considered valid only in the hot postshock slab and in the corona (21). There, the thermal conduction together with the radiative losses from optically thin plasma plays a significant role in the energy budget. In particular, the intense radiative cooling at the base of the slab robs the postshock plasma of pressure support, causing the material above the cooled layer to collapse. As a result, the shock position can vary in time (25). The thermal conduction acts as an additional cooling mechanism of the hot slab, draining energy from the shock-heated plasma to the chromosphere, and partially limits the growth of thermal instabilities (25).
On the contrary, the cold and dense material of the stream and that of the chromosphere are most likely optically thick. As a result, the radiative transfer is expected to play a significant role in the energy budget, whereas the thermal conduction should be negligible. The main effects are expected in the unshocked accretion column where the downfalling material can be radiatively heated to temperatures up to 10 5 K (21). Also, the optically thick material of the chromosphere and/or of the unshocked stream located along the line of sight (LoS) is expected to partially absorb the x-ray emission arising from the hot postshock slab. Note that we do not consider the effects of radiative transfer on the dynamics and energetics of the system, but we account for the absorption in the synthesis of x-ray emission, as described in the “Synthesis of the x-ray emission and comparison with astrophysical objects” section. For this reason, our modeling is not entirely self-consistent. Nevertheless, we expect that the evolution of the hot (T > 1 MK) postshock plasma is accurately modeled by radiative cooling.
Both the shocked core plasma and the shell are simultaneously observed in the recorded laboratory plasma emissivities. The reverse shock front and its temporal evolution, propagating up the stream at
14 ± 3 km/s, are clearly seen in the streaked visible emission (see Materials and Methods) of the laboratory plasma (Fig. 2A): The reverse shock front identified in the density maps (Fig. 2A, red points) corresponds closely to the edge of the bright emitting, core postshock region that expands toward the incoming stream. In the same emission map, we can also clearly identify the shell, of reduced brightness, with its expansion front in the density maps (Fig. 1C, yellow points). Similarly, the x-ray laboratory emission originating near the obstacle surface (Fig. 2B), analyzed by our nonsteady model (37) and detailed in the Supplementary Information, displays features characteristic of two distinct plasma components. Here, the appearance of intense He-series lines (from emitting He-like ionized F ions) is the witness of a plasma component having a low temperature (0.6 ± 0.1 MK) at a density that corresponds well to the core density observed in Fig. 1C. The simultaneous observation of a strong Lyα line (from emitting H-like ionized F ions) attests the presence of a higher electron temperature plasma, that is, at 3.7 MK, analyzed using the atomic code FLYCHK (41). This synthetic radiation, derived from the ratio between the Heβ and Lyα line intensities, has a volume and a density consistent with those of the measured shell plasma (Fig. 1C). Note that both shell and core temperatures derived this way are also well consistent with the laboratory simulation, as detailed in the Supplementary Information.
Title: Tidal Disruption of a Main-Sequence Star by an Intermediate-Mass Black Hole: A Bright Decade
10 yr long super-Eddington accretion phase. The photospheric emission of the outflow ejected during this phase dominates the observable radiation and peaks in the UV/optical bands with a luminosity of 10^42 erg/s. After the accretion rate drops below the Eddington rate, the bolometric luminosity follows the conventional t^ <-5/3>power-law decay, and X-rays from the inner accretion disk start to be seen. Modeling the newly reported IMBH tidal disruption event candidate 3XMM J2150-0551, we find a general consistency between the data and predictions. The search for these luminous, long-term events in GCs and nearby dwarf galaxies could unveil the IMBH population.
Title: Detection of a Cool, Accretion Shock-Generated X-ray Plasma in EX Lupi During the 2008 Optical Eruption
0.4 keV plasma component, as expected for accretion shocks on low-mass, pre-main sequence stars. From 2008 March through October, this cool plasma component appears to fade as EX Lupi returns to its quiescent level in the optical, consistent with a decrease in the overall emission measure of accretion shock-generated plasma. The overall small increase of the X-ray flux during the optical outburst of EX Lupi is similar to what was observed in previous X-ray observations of the 2005 optical outburst of the EX Lupi-type star V1118 Ori but contrasts with the large increase of the X-ray flux from the erupting young star V1647 Ori during its 2003 and 2008 optical outbursts.