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The Evolution of Coronal X-ray Emission The Sun in Time Rob Jeffries Keele University NASA GSFC The Sun has been around for 4.6 Gyr. After 10-20 Myr as a pre-main sequence star it settled onto the ZAMS as a hydrogen burning star. Rocks were present in the solar system at 1 Myr, the giant planets formed within 10 Myr and the Earth assumed its final mass on timescales of 20-30 Myr. Since then the Sun has increased in luminosity by 30%, but conversely its EUV and X-ray coronal emission has declined by many orders of magnitude. Ultimately this is driven by angular momentum loss and the connection between the amplification and emergence of B-fields and rotation. It is the dissipation of magnetic energy that heats the corona and Bfields that confine it. In this lecture I will be looking at what we know of the past history of solar coronal emission. 1 Manifestations of the Sun s magnetic field The Sun is a magnetically active star. The magnetic field is produced by an interaction between rotation, differential rotation and convective motions. This magnetic dynamo amplifies magnetic fields, which become buoyant and emerge from the photosphere. These fields frequently give rise to dark sunpots, where the field has suppressed the upwelling of hotter material from beneath. The Sun of course looks different at X-ray wavelengths. Loops of magnetic field constrain plasmas at millions of degrees. The footpoints of these loops cover at most a few percent of the solar surface and the plasma is confined to heights only a fraction of the solar radius. 2 We know that the solar activity varies in a cyclic way, such that things like the X-ray luminosity and the rate at which flares and coronal mass ejections occur go through order of magnitude changes on timescales of 11 Myr, whilst the optical output varies by just 0.1%. 3 X-ray Irradiance variation: GOES 1-8 Angstrom Strong & Saba 2009, Adv. Sp. Res. 43, 756 On the other hand the X-ray luminosity of the Sun varies by more than an order of magnitude over the same solar cycle. 4 Planetary interactions, with winds and radiation An interesting question is to ask what did the Sun look like in the past? What was the overall level of its high energy irradiance? By how much did it vary? How frequently did it flare? Etc. Why is this important? Aside from the intrinsic interest of understanding how the mechanisms of generating and dissipating magnetic field operate at high rotation rates there are the important topics of: How the solar corona flares, coronal mass ejections etc. interact(ed) with planetary systems. The properties of the early Sun in this regard may be responsible for the way the atmospheres of planets in the inner solar system have turned out (i.e. photoionisation/photodissociation can change chemistry, X-ray/EUV radiation can heat upper atmospheres leading to excessive escape of species; a stronger solar wind can strip ions from atmospheres unless they are protected by magnetic fields). This may account for the dry atmosphere of Venus and the thin atmosphere and removal of water from Mars. None of which I am able (or qualified) to discuss any further here. At even earlier ages, the impact of high energy radiation on a circumstellar disk may have been important in driving mass loss and mass accretion and for certain chemistry. 5 The solar magnetic bubble 100 AU NASA JPL On a larger scale, the solar system sits within a magnetic bubble blown out by the solar wind. To some extent this bubble shields the solar system from much of the cosmic ray flux. If the Sun s magnetism was stronger in the past, this shield would presumably have been more effective. 6 How do we infer the past properties of the Sun? The rotation-age connection The age-dependence of X-ray luminosity and the high-energy spectrum Variability, cycles and flares Magnetic and coronal structures The X-ray activity of pre main sequence Suns Here is what I will speak about. First, how do we infer the past properties of the solar corona? One way is to look for archaeological evidence in the form of the composition of ancient meteorites or surface composition of the lunar surface or indirect evidence from planetary atmosphere histories. However I will not talk about these. I will talk about using solar analogues stars of a similar mass to the Sun but with younger and older ages. Detailed observations of these can reveal how the coronal activity of a Sun-like star has changed with age including luminosity, coronal temperatures, flare rate, activity cycles and coronal structuring. We will see that at the very youngest ages, the Sun s history is not completely determined and indeed when we talk about pre-main sequence stars, some of the X-ray activity that is seen does not entirely have any analogy in the present-day Sun; and for that reason I treat these objects separately. 7 What changes about the Sun once it has reached the ZAMS after about 20 Myr? Stars are spun down by a magnetised wind Spin down is greatest for faster rotators This leads to convergence and a single-valued rotation-age connection for stars older than 500 Myr. age-0.6 0.1 1.0 10 Ayres 1997, JGR, 102, 1641 Once the Sun has reached the ZAMS then the overall progress of its evolution is rather stately, governed by the nuclear fusion timescale. The Sun is probably about 30% more luminous than it was on the ZAMS, unless there has been unexpectedly large mass loss in the past (faint young Sun paradox). The contraction along the Hayashi track will leave the ZAMS Sun rotating much faster than it does today. Exactly how fast is somewhat moot, but rotation rates between 5 and 100 times as fast are seen in young Suns in open clusters of known age. The Sun get s to where it is today by spinning down via a magnetised wind. The ionised outflow locks to open magnetic field lines until it decouples at some distance well away from the solar surface. This process removes angular momentum. The spin down rate depends on the strength of the stellar magnetic field and the mass loss rate, such that an early dispersion of rotation rates will converge to a single relationship if B increases with omega, as suggested by dynamo models. In fact the relationship observed is reasonably well described by B propto omega. This means we cannot for sure say what the rotation rate of the Sun was at ages <500 Myr where this divergence takes place. 8 Log Period More data! Solar analogues Almost a single-valued PeriodAge relation Log Age Big spreads here Convergence here Young clusters A large scatter at young ages Converges >500 Myr Bouvier 2008, A&A, 489, L53 Age We can look at this in a little more detail. At older ages a single power law relation fits a number of solar analogues, but there are order of magnitude spreads at younger ages. Current models of rotational evolution need to postulate an initial spread in angular momentum plus a variation in the amount of time objects spend magnetically coupled to the inner parts of their protoplanetary discs during the first 10 Myr. 9 1978-1981 HEAO-2 Einstein 0.1-4 keV imaging (60 ) proportional counter 1983-1986 EXOSAT 0.05-2 keV imaging (18 ), 150 keV proportional counter 1990-1999 ROSAT 0.1-2.4 keV imaging (15 ) proportional counter 1992-2001 EUVE 0.01-0.2 keV spectroscopy (R=200) 1999- Chandra 0.2-10 keV CCD imaging (0.5 ) and spectroscopy (R~1000) 1999- XMM-Newton 0.1-10 keV CCD imaging (5 ) and spectroscopy (R=800) Now to observe the changing X-ray properties of solar analogue, we need X-ray telescopes. X-rays cannot penetrate the atmosphere, so a series of increasingly sophisticated satellites have been flown. Einstein was the first to detect significant numbers of coronal sources from normal stars. Both it and the European follow-up EXOSAT had very limited spatial and spectral resolution. ROSAT offered reasonable imaging, a much larger telescope area and a resolution capable of broadly characterising coronal temperatures using proportional counter spectroscopy (X-ray colours). Grating spectrographs were flown on several subsequent satellites. EUVE was a relatively small EUV telescope. Chandra offers superb spatial resolution and excellent low energy spectroscopy. XMM-Newton has a larger collecting area, lower spatial resolution, but better spectral resolution at high energies. 10 The rotation-activity connection X-ray emission depends on the square of the rotation rate Pallavicini et al. 1981, ApJ, 248, 279 The next piece in the jigsaw is to examine the relationship between rotation rate and X-ray activity. Very early on it was realised that there was a close connection between rotation and activity, both in the corona via X-ray emission (observed by Einstein) and in the chromosphere observed at UV (Ca H+K and IUE). The dependence is roughly a power law with index 2. 11 X-ray emission depends on the square of the rotation rate In and Why? dynamo models For stars of a given spectral type B Durney & Robinson 1982, ApJ, 253, 290 And the energy density of these fields U B2 But exactly how this energy density is converted into X-rays is one of the great unsolved problems in Astrophysics! The basic argument is a simple back of the envelope thing. Various analytic dynamo models suggest that the B-field emerging at the surface of a star will be propto the angular velocity or rotation rate. The energy density in these fields is proportional to B^2, hence the squared dependence on rotation rate, The missing step is to understand how that magnetic energy is dissipated and converted into X-rays. 12 More recent data a clear dependence on rotation but much scatter Lx/Lbol 2 Includes stars of all spectral types Sun Pizzolato et al. 2003, A&A, 397, 147 More recent data has firmed up this conclusion but also revealed a more complex picture. At relatively slow rotation rates there is a decline with decreasing rotation rate. The Sun more-or-less fits in with this relationship. At periods < 5 days there appears to be a turnover or saturation in the fraction of luminosity emitted at X-ray wavelengths. There is a fair bit of scatter in these relationships and it appears that this is because we are mixing in stars of all spectral types. In fact, dynamo theory suggests that not only rotation rate is important, but that dynamo efficiency will be proportional to the time it takes a convective cell to rise and fall, the so-called convective turnover time. Ie.. Lx propto (turnover time/rotation period)^2, the square of the inverse Rossby number. 13 Log (Lx/Lbol) Saturation Sun Log (Rossby Number) Unification using the Rossby Number The ratio of period to convective turnover time Jeffries et al. 2011, MNRAS, 411, 2099 Plotting X-ray activity vs Rossby number cleans up the diagram considerably, but the saturation of magnetic activity at fast rotation rates (or deep convection zones with long turnover times) remains. This saturation is now thought to be a property of the dynamo mechanism itself, a Lorentz feedback on the flows that generate the fields. That stars of very different spectral types (and even PMS stars/giants) appear to saturate at the same Rossby number is a compelling argument for this. 14 Important Plot! X-ray luminosity vs age for solar-type stars Saturation Lx/Lbol=10-3 Spread caused by rotation Spread caused by variability So here is the first very important plot. A compilation of the X-ray luminosities for G0-G5 stars of known age (either because they are members of open clusters, they have been associated with kinematic groups of known age or their age has been estimated from the HR diagram (for the older stars). Similar spectral types remove the concern about convective turnover time dependence. For the clusters, the bars shows the 20th to 80th quartile. For the field stars the bar links maximum and minimum observed fluxes where objects have been observed multiple times. Points to note: The Sun is 2-3 orders of magnitude fainter than 100 Myr old ZAMS stars. It also varies in flux by >1 order of magnitude as does alpha Cen A at a similar age. Younger Suns are much more luminous. However multiple observations do not betray any large variability, yet there is a spread of luminosities within a cluster, why is this? Well it is actually because of a spread in rotation rates at young ages. In fact the young, coronally active stars hardly vary at all. 15 X-ray variability LX (1995) Sun varies by factor of 20 on 11 year cycle 7 year gap Active stars vary by much less than a factor of 2 on 10 year timescales NGC 2547 The young sun (<0.6 Gyr) did not exhibit large amplitude X-ray variability on 1-10 year timescales. Age 35 Myr LX (2002) Jeffries et al. 2006, MNRAS, 367, 781 This shows an example (there are others) of stars in a cluster being measured twice and hardly any even showing factor of 2 variations on 10 year timescales. Whilst it is possible that they just have very long cycles >>10 years, the average luminosities as they stand fit in well with the rotation activity relationship. If stars scattered by a magnitude on longer timescales we would expect to see a similar scatter in rotation-activity relationships, which is not observed. 16 Factor 2 X-ray variability? Guedel 2004, ARAA There is some tentative evidence of X-ray, and better evidence for starspot, cycles on 0.1 Gyr solar analogues. But the X-ray variation is only a factor of 2. X-ray maximum *probably* coincides with min optical brightness and hence max spottiness as expected. A difficulty here is that measurements are taken with different instruments with different responses difficult to normalise. 17 X-ray variability But at some point (older than the Hyades at 600 Myr), order of magnitude (possibly cyclic) variability occurs. In the past the Sun was much more luminous but much less variable on long timescales. G2V Age ~5 Gyr K1V Robrade et al. 2012, A&A, 543, A84 The long mission lives of XMM and Chandra have allowed long term monitoring of a few solar type stars. It seems that larger X-ray variations are present in older stars like the Sun. There are now a couple of examples where the variation appears to approach that seen in the Sun. We can conclude that whilst the ZAMS sun had a much higher X-ray luminosity than at present, it did not show large amplitude solar-type X-ray cycles on timescales similar to the solar cycle (or shorter). 18 XMM CCD Spectra Normalised Countrate X-ray spectra drop in flux and get softer with age 0.1 Gyr 0.3 Gyr 1.6 Gyr Photon energy (keV) Telleschi et al. 2005, ApJ, 622, 653 We can also examine how the energy distribution, reflecting the temperature distribution of the dominant plasma, changes with time. The CCD spectra offered by XMM (and Chandra) allow the gross spectral distributions to be compared rather easily. It is obvious that as the solar analogues get older, their luminosity drops but their spectra get softer. I.e.the reduction in higher energy X-rays is quicker, suggesting that the dominant coronal temperatures get cooler. 19 Age (Gyr) XMM Grating Spectra Ne X Fe XVII Different species probe different plasma temperatures O VIII 0.1 e.g. > 10 MK - Ne X < 5 MK - Fe XVII 0.1 O VII 0.3 Fe XVII 0.3 0.75 Must have multitemperature coronae 1.6 Telleschi et al. 2005, ApJ, 622, 653 However, it s a bit more complex than a simple coronal cooling. The grating spectra of the same analogues reveal a more complex picture. Different species with different excitation energies are sensitive to plasma at different temperatures. So, whereas the overall picture is a lower average temperature with age, it is clear from the presence of lines sensitive to very different temperatures, that the coronal plasma is at multiple temperatures or has a continuous differential emission measure (DEM) distribution. 20 Differential Emission Measure 0.1 Gyr EM-weighted Average T 0.3 Gyr 1.6 Gyr X-ray Luminosity Telleschi et al. 2005, ApJ, 622, 653 The grating spectra can be modelled with a continuous DEM. The results are somewhat non-unique, but consistent treatment of the solar analogues tells the expected story that there is a reduction in the amount of high temperature plasma as a star gets older. If we average over this DEM to calculate some sort of average temperature then there is a strong dependence of X-ray luminosity on temperature. The hotter coronae are much more luminous. The young Sun is about 5 times hotter and 1000 times more luminous. 21 Extend into the EUV/FUV Flux (age) at 0.1 MK, at 10 MK, 2 High Temperature plasma diagnostics decay more rapidly Ribas et al. 2005, ApJ, 622, 680 This work has now been extended into the UV and EUV and an interesting picture emerges. The higher temperature plasma diagnostics and harder flux decays more rapidly. This has important implications for the irradiation of planetary systems, because different wavelengths of radiation have different ionisation, heating and penetrative capabilities. 22 From Guedel 2007, LRSP, 4, 3 To put some flesh on this, here is a table with the numbers. So at 0.1 Gyr, whilst FUV radiation is enhanced by a factor 25 over the present day average sun, the X-ray emission is enhanced by a factor of about 1000, though the exact number depends on the rotation rate of the sun at that time, which is as-yet unknown. 23 X-ray spectra INCREASE in flux and get HARDER at younger ages possible explanations? 1. Higher coverage by high pressure or large coronal loops in the past? T = 1400 (p L)1/3 Lx n2 L3Q(T) LT3.7 2. Superposition of power law flare distribution. Young stars have more energetic and hotter flares? So what are the possible explanations for all this. Well, the dynamo efficiency increases with omega and this supplies the raw energy for the process, but how is this manifested. Well perhaps the surfaces of these stars are covered by more active regions and perhaps the magnetic loops in those active regions are filled with plasma at high pressures. As the emission measure scales as density squared, then for loops in equilibrium there is a natural connection between luminosity and temperature for loops of a given size. An alternative is that there is no such thing as a static loop and that the coronae of these stars consists of a multitude of flaring loops, and this works because young stars have more and more energetic, hotter flares than the Sun. 24 Simple loop models Quasi-static balance between heating and radiative/conductive cooling *RTV scaling law T = 1400 (p L)1/3 p T3 L-1 * Rosner, Tucker, Vaiana 1978, ApJ, 220, 643 N loops, area A, length L EM = ne2 NAL Filling factor f f = 2NA/2 R2 Pressure p p = 2nekT p EM0.5 T (f L)-0.5 Looking first at static loop models. The very famous RTV paper establishes a scaling law between pressure, temperature and loop length. Note this is non-unique. A given temperature can be produced by a locus of pL. However, if we have an array of N loops whose footpoints cover a fraction f of the visible surface, this defines another relationship between p, T and L that can break the degeneracy. 25 RTV Sun T=3x106 K EM=3x1049 cm-3 f=0.01 f=1 Active Regions So taking typical solar coronal temperatures and emission measures we see that the Sun s coronal emission could be dominated by a few active regions. Indeed, plasma pressures in active region loops are measured to be 1-10 dynes/cm^2 and have lengths 0.01-0.1. This is consistent with the solution above which has these regions occupying a few percent of the surface. So even though it is seriously doubtful that ideal static loops exist, the method seems to have some merit. 26 f=0.01 EK Dra T=107K EM=3x1052 cm-3 f=1 Active Regions RTV Applying this to a young solar analogue we see that even if f=1, complete coverage by solar-type active regions would fail by more than an order of magnitude to explain the observed emission measure. The loops would have to have much higher pressures and/or larger loop heights/volume to produce the necessary emission. So this is not just scaled up solar activity. 27 Guedel et al. 1995, A&A, 301, 201 EK Dra Rotational modulation makes it unlikely that f=1 or that loop heights are very large However, the fact that rotational modulation at the level of 50% can often be found makes it very unlikely that f can be approaching 1 and also constrains plasma scale heights to be less than a solar radius. But given that the hydrostatic pressure scale height is also less than a solar radius and that emissivity scale height is half of this, the latter is hardly surprising. Either way it implies that the coronal pressures must be much larger than on the sun. 28 Solar Max Solar Min Flaring 106 107 Temperature (K) Compare Sun as a star with EK Dra Very little high temperature plasma, found mainly in flares Argiroffi et al. 2009, A&A, 488, 1069 A clue is to compare the DEM of the quiet and flaring sun. The only time that plasma temperatures in excess of 1e7K are seen is at solar maximum, especially during solar flares, So it is natural to ask whether the increased luminosity and DEM of young solar analogues, which is largely on the hard X-ray side, is due to an increased rate/size of X-ray flaring. It has been shown that a large superposed population of randomly flaring loops that are then allowed to cool will approximately reproduce the observed DEM of very active solar analogues. Is there any evidence that the corona is predominantly heated by flares? 29 Flare rates N(>E) (day-1) 0.1 Gyr C-rate E (ergs) N(>E) (day-1) 0.75 Gyr days dn/dE = KE- Sun E (ergs) Audard et al. 2000, ApJ, 541, 396 In the sun it has been observed that the energy spectrum of hard X-ray flares is a power law of index about 2. If dN/de = kE^-2 -> E = int E E^-2 dE between a high and a low cut off. If the lower cut-off is very small, then [ln E]^Eup_0 = infinity, and nanoflares could heat the entire corona. In fact the energy spectrum of young solar analogues has a slope of 2.3 meaning that even a finite low E cut-off can result in the entire luminosity of the star being produced by flares. The slopes are very similar in stars of differing age, but the normalisations are different by orders of magnitude. In the Sun, the largest flares of 1e32 erg occur of order once per month. In a 0.1 Gyr solar analogue this is below the threshold of what can be seen, but extrapolation suggests >10 per day! Flares that are 1000 times more powerful occur once per month. 30 Young, active Suns flare more, and flare bigger Even observed flares explain a good fraction of Lx Audard et al. 2000, ApJ, 541, 396 Flare rate correlates with Lx Even the flares that are observed can account for 10% of the observed Lx, without any extrapolation to lower energies. However, the rate at which large flares occur correlates well with Lx, suggesting that there is a deeper connection between flaring and coronal heating. 31 Other evidence for the importance of flare heating Correlation of U-band flare energy with X-ray luminosity Doyle & Butler 1985, Nature, 313, 378 Correlation between non-thermal radio- and X-ray luminosities Guedel & Benz 1993, ApJ, 405, L63 X-ray luminosity Solar flares Young solar analogues Radio luminosity These are not new ideas. As long ago as 1985, our esteemed organiser noted that the amount of energy in U-band flares was well correlated with Lx. There is also a correlation covering many orders of magnitude between nonthermal radio emission (produced by flare acceleration events) and Lx. But why are there more and more energetic flares? Possibly the packing together of more magnetic structures combined with some jostling/differential rotation and possibly larger plasma volumes leads to more frequent flaring and when those flares occur there is more magnetic energy available to dissipate. This is a simplistic picture and we need to know something about the structure of the magnetic field and corona to see whether this is consistent is there any evidence for larger magnetic filling factors, larger coronal heights etc? 32 Coronal/Magnetic Structure in time: (Adolescents are spotty!) Typical light curve amplitudes 0.03-0.1 mag 3-10% spot coverage? c.f. only 0.1% in the Sun EK Dra Start off with the photosphere. Young stars show modulated light curves, normally interpreted in terms of asymmetric spot coverage. If this were modelled using a single spot we would deduce spot areas of 3-10% (c.f. 0.1% in the Sun). 33 Spots are probably much more extensive Big change in unspotted level Messina et al. 2002, A&A, 393, 225 10% more spots? Light curves only measure axi-asymmetric component However, this is just a minimum coverage just the asymmetric component. Long term light curves shows evidence of activity cycles, but also shows a 0.1 mag overall modulation so there are at least 10% of unmodulated spots. This is still a lower limit. 34 Coronal/Magnetic Structure in time: (Adolescents are spotty!) Doppler Imaging HD171488 P=1.34d HII 314 P=1.47d EK Dra P=2.7d Strassmeier et al. 2003 Rice & Strassmeier 2001 Strassmeier & Rice 1998 Doppler imaging shows spot coverage fractions of 10-20%, but again may be a lower limit because if small spots are present they will be unresolved. The resolution of doppler imaging is currently about 10 degrees which is about 10 times bigger than sunspots. There is some evidence for a polar concentration of spots in some stars, very un-solar behaviour. 35 TiO band modelling of 0.1 Gyr solar analogue O Neal et al. 2004, AJ, 128, 1802 fspot Rotation phase fspot ~0.4 from TiO bands (c.f. <0.01 on Sun) Finally, TiO band modelling of EK Dra suggests that actually there is a 40% spot coverage of spots that are 1000-1500K cooler than the unspotted photosphere. This is a massive coverage and suggests the corona may look quite different too, with loop footpoints likely to cover a much larger surface area of the star than in the solar case. 36 Alpha CrB A0+G5V (0.2-0.5 Gyr) Mapped via its eclipses Guedel et al. 2003, A&A, 403, 155 But how to resolve the corona? Basically we can t. We can try indirect imaging utilising eclipses, but there are hardly any suitable systems. Here is one example where the A-star companion is almost certainly X-ray dark so acts as an occulting disk. Modelling of the light curve (which is a nonunique inversion) suggests that the corona around the G-star (age 0.2-0.5 Gyr) is relatively compact and has a relatively low filling factor (but still large by solar standards). 37 (r) (i) Density sensitive line ratios (f) e.g. He-like triplets Porquet et al. 2010, SSRv, 157, 103 Another way of deducing coronal size is to look at coronal densities. To produce a given emission measure (n_e^2 V), then a high density corona would have a smaller volume. There are a variety of spectroscopic techniques; perhaps the best is to use Helium-like triplets. One transition is forbidden, one allowed and there is an intercombination level that can be partly filled by collisional interactions. The ratio between the forbidden and intercombination lines is density sensitive, down to a low-density limit where collisional excitation becomes unimportant and the ratio of both of these to the allowed transition is temperature sensitive. Problems: Is the plasma optically thin? Are there faint, blended lines? Is the background UV field negligible? Any measurement of density is weighted (by the emissivity) towards higher density regions. 38 r f i Typical derived densities 1010-1012 cm-3 Ness et al. 2002, A&A, 394, 911 These measurements are possible at the grating resolution of Chandra and XMM. Typical derived densities are 10^10 to 10^12 cm^-3. There are various possible problems here, both in the atomic physics, dealing with blends with the relevant lines and possible optical depth effects. 39 EK Dra T=107K f=0.01 EM=3x1052 cm-3 f=1 Range of implied coronal pressures Active Regions RTV Coronal densities in active stars suggest N_e of anywhere between 10^10 and 10^12 depending on which diagnostics are used. At temperatures of 10^7K this means pressures of 10-1000 times higher than solar active regions. This suggests that the visible corona (that which contributes to the density diagnostics) is less than a solar radius if the plasma is anywhere near equilibrium and also that f<<1 if scale heights are indeed less than a solar radius. This is self-consistent with the presence of rotational modulation of X-ray emission in some stars. 40 Guedel et al. 1995, A&A, 301, 201 EK Dra Rotational modulation makes it unlikely that f=1 or that loop heights are very large However the fact that rotational modulation at the level of 50% can be found makes it very unlikely that f can be approaching 1 and also constrains plasma scale heights to be less than a solar radius. But given that the pressure scale height is less than a solar radius and that emissivity scale height is half of this, the latter is hardly surprising. 41 AB Dor, rapid rotating young sun, 0.5d period Co-rotating Halpha prominences held in Bloops far from stellar surface Collier Cameron & Robinson 1989, MNRAS, 238, 657 So, we have the picture of a compact, high pressure corona. However, there are other indications for much more extended structures, which though they may not contribute a huge emission measure because of the n_e^2 dependence, they could yet contain much magnetic energy (propto B^2 x volume). For instance in many rapidly rotating stars there are indications of co-rotating cool prominences situated at or around the Keplerian co-rotation radius (>1 Rstar above surface). It seems extremely likely that these must be magnetically confined, and indeed the presence of instabilities leading to centrifugally compressed, cooling plasma is predicted if loop apexes exceed the Keplerian co-rotation radius. 42 The Pre Main Sequence Sun <10 Myr Discs, accretion, outflows At early times <10 Myr, there are other phenomena to consider. These include that a large fraction of such stars possess circumstellar discs from which the star accretes and which may also drive strong outflows along their rotation axes or from the surfaces of the discs themselves. 43 The plot on the right shows a general picture of how circumstellar material and accretion develop. Class 0 and Class I protostars are enshrouded by an optically thick envelope that is bright only at mid to far infrared wavelengths. As the envelope is dispersed or accreted, the T-Tauri star is revealed, surrounded by a geometrically thin, but optically thick disk from which it may be accreting. Some time later, on timescales of a few Myr, most stars lose these discs, which is signalled by the loss of emission-line accretion signatures and the disappearance of any infrared excess. In stars older than 10 Myr, any circumstellar material is opticallly thin and likely to be second generation material generated by a forming planetary system. In addition, the structure of a PMS star is very different. In this PMS stage there are of course no H-burning reactions, though there will be some minor contribution from early deuterium burning. In addition, there is thought to be no radiative core in a 1 solar mass star before about 10 Myr as it descends the fully convective Hayashi track. As such one might expect a different Type of dynamo to operate compared with the Sun (where the dyanmo is thought to operate at the interface between radiative core and convective envelope the tachocline), if indeed one operates at all. 44 Orion Nebula Cluster Chandra Age 2 Myr Wolk et al. 2005, ApJS, 160, 423 So what do such objects look like in X-rays? This is an image of the Orion Nebula Cluster (age 2 Myr) taken with the Chandra satellite at energies of 0.5-8keV range. There are more than a 1000 X-ray sources coinciding with PMS stars. Highlighted here are those which are solar analogues with spectral types of K5-K7 at this age. In general, protostars of classes I, II and III are all found to be strong X-ray emitters. 45 Accretion shocks? X-ray emitting loops A number of sites for this X-ray emission have been suggested. These include conventional coronal X-ray loops heated by similar processes to ZAMS stars; or emission from a shock region where magnetically channelled material falls from the disk onto the stellar surface. The same process is thought to be important in regulating the rotation of PMS stars and enabling them to shed large amounts of angular momentum through the disk. 46 Source are highly variable and flare frequently Plasma temperatures 107 to 5x107 K 13 days Suggests a coronal origin Wolk et al. 2005, ApJS, 160, 423 The X-ray emission from these sources is highly variable but shows clear signatures of flares (impulsive rise, longer decay). The inferred plasma temperatures are as hot or a bit hotter than the hottest ZAMS stars. Both of these suggest a coronal origin for most of the emission, though there has been evidence that at least part of the soft X-ray emission <0.5keV could be produced in accretion shocks. 47 Lx/Lbol = 10-3 Updating the diagram we used earlier, we see that on average these young PMS stars are more X-ray luminous even than ZAMS stars. However, they are bolometrically more luminous too. In fact most PMS stars have Lx/Lbol of around 1e-4 to 1e-3 which is the saturation limit in older stars. There is however more than an order of magnitude variation in the X-ray luminosity of PMS stars. Why? Well it is not because of long period variability like the Sun so is it like the young clusters that there are a range of rotation periods? The periods of PMS stars are in the range 1-10d, so might this account for it? 48 Little correlation with rotation or Rossby number large scatter Locus occupied by MS stars ONC PMS stars Preibisch et al. 2005, ApJS, 160, 401 Well no. In fact there is almost no correlation with rotation found. The reason is likely that because of their deep convection zones almost all PMS stars are in the saturated part of the activity-Rossby number plot. But, the scatter is *much* larger than is seen among fast rotating ZAMS stars, and much larger than any plausible observational uncertainties, so what could account for this? 49 Accreting stars are less active Accretors Nonaccretors Preibisch et al. 2005, ApJS, 160, 401 One observed fact is that stars that are actively accreting are in general less coronally active than stars without disks and accretion. Why is this? (1) It could be due to absorption of the X-rays by disk material or accretion veils, but the signature of this in the spectra is not always obvious. Another possibility is that the mass-loading of a fraction of the magnetic field lines makes it impossible to heat these up to coronal temperatures. This is still very much an area of active research. 50 Strong flares E ~ 1035-1036 ergs once per week dn/dE E-1.7 Wolk et al. 2005, ApJS, 160, 423 Finally it s worth looking at the flaring statistics of PMS stars. Long observations of star forming regions can gain many time series on solarmass stars. The results are quite similar to active ZAMS stars. The power law slope is close to 2 and there are some truly gigantic flares. On a typical solar-type PMS stars there is at least one flare with an energy exceeding 1e35 (1000 times as large as the largest flares on the Sun) occurring every week. The modelling of the flare light curves suggests that in some cases these occur in structures that are several stellar radii in size, perhaps in structures linking the star to its disk. On the other hand rotational modulation in many stars suggests that relatively compact emission is also common, though not with filling factors approaching unity. 51 X-ray/EUV irradiation of the disc surface Photo-ionisation, disc heating, photochemistry Typical X-ray luminosities will dominate photo-ionisation of discs out beyond 100 au, exceeding the influence of cosmic rays. This photoionisation can couple the disc to magnetic fields and it becomes unstable to magneto-rotational instabilities that act as a form of viscosity that drives mass accretion. Heating of the upper layers can also result in photoevaporation and may determine the disk lifetime. Finally the high energy photons can initiate certain chemical reactions. 52 Summary The Young Solar Corona: 1. Was much more luminous (x 1000) 2. Flared more frequently and more violently (x1000) 3. Was hotter (x5) 4. Had a higher density (x10) 5. May have been dominated by flaring loops 6. Was more densely packed with magnetic structures In summary, this is what we know reasonably conclusively about the young Sun. We do not however know in detail (to the nearest factor of a few) just how luminous it may have been when younger than a few hundred Myr, because of the ambiguity about the early solar rotation rate. Also a major uncertainty is exactly what the corona looked like. There is evidence for both compact high density loops, but also for more extended loops beyond a stellar radius. How the B-field is organised is still somewhat uncertain. The filling factor of magnetic field structures must have been much higher and it seems quite likely that this might account for enhanced rates and sizes of flaring activity in the corona. 53