Multi-wavelength observations of variability characterizing magnetic activity in late-type stars.

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May 19, 2017

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Multi-wavelength observations of

variability characterizing magnetic activity in late-type stars

D. Pizzocaro

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General outlook

This thesis consists of the presentation of two diﬀerent works, whose common factor is their relevance to the ﬁeld of the magnetic activity of late-type stars.

The presence of magnetic fields on the surface and in the outer atmosphere of the Sun and other late-type stars determines the formation of small scale structures which make the atmosphere inhomogenous, and causes, through the so-called “magnetic heating processes”, excess heating of the atmospheric plasma which consequent excess radiation. This is what is referred to as magnetic activity. Magnetic processes also cause non-thermal radiation, due to accelerated particles (in flares). The study of magnetic activity gives es- sential information on the properties of the stellar magnetic fields, on their strength, geometry and, eventually, on their origin through the internal dy- namo mechanism. Apart from the comprehension of the physical processes occurring in the atmosphere of stars and in stellar interiors, magnetic activity has, together with spectroscopy, a crucial role in determining the physical and chemical characteristics of the cirumstellar medium, because of the associated high-energy emission, which plays a key role in processing the circumstellar material. Together with stellar winds and energetic phenomena such as the Coronal Mass Ejections, the magnetic activity may have a strong influence on the atmospheres of planets and influences the habitabil- ity of planetary systems. The effects of the magnetic-activity related high energy emission is particular relevant during the first phases of the life of a late-type stars, which are known as the “Young Stellar Object” (YSO) phase.

In the first work I present, I analyse the activity of a sample of late-type main-sequence stars observed at optical wavelengths by the Kepler mission and, in X-rays, by the XMM-Newton mission, and I characterize the rela- tion between magnetic activity and the stellar rotation, which is a proxy for the internal dynamo mechanism of stars. In the second work, I present the discovery of strong magnetic activity, in the form of an X-ray flare, from a very young (Class I) YSO, ISO-Oph 85; this latter work is focused on the characterization of this transient X-ray emission and on the classification of

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6 Chapter 1. General outlook the evolutionary stage of ISO-Oph 85.

In both the works I present, the study of X-ray activity plays a major role. The X-ray data I use are taken from the XMM-Newton observatory.

Since variability is an intrinsic feature of magnetic activity, I take advantage, for the time-resolved analsyis of the X-ray data, of the algorithms and prod- ucts provided by the EXTraS Project, which is aimed at the characterization of variability in the database of XMM-Newton. In the ﬁrst one of the two works presented, the observations performed by the Kepler mission have a crucial role.

In the introductory Chapter 2, I present some aspects of the magnetic activity that are useful to the comprehension of the topics discussed in this thesis. Magnetic activity has been discovered on the Sun: sunspots are the ﬁrst observed feature due to inhomogeneities in the geometry and emission properties of the solar atmosphere caused by the presence of magnetic ﬁelds.

The Sun represents the “home lab” in which activity can be studied closely, and the paradigm of our knowledge of magnetic activity, which has been then applied to the stars searching for differencies and similarities. A description of the activity-related phenomena that can be observed in the solar atmosphere, and their discovery, is given. Then I present the theory concerning the origins of the magnetic field of the Sun through the dynamo mechanism, which can be extendend, with some limitations, to late-type stars. I also discuss the case of two peculiar classes of stars (with fully convective and fully radiative interiors), and how their activity differs from solar-like stars.

In the next section I describe the different manifestations of activity that can be observed on solar-like main-sequence stars, and the relation between activity and rotation. Lastly, I describe briefly the characteristics of YSOs and the peculiarity of their activity, with particular attention to the flaring activity.

XMM-Newton, the EXTraS Project and Kepler are described in Chapter 3. The works which are th subject of this thesis are presented in Chapter 4 and Chapter 5.

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Chapter 2

The magnetic activity of late-type Stars

2.1 The Sun and the discovery of magnetic activity

The Sun has been defined “the Rosetta stone of the Astrophysics” (Landi Degl’Innocenti, 2008). This definition is justified by the fact that the observation (and the understanding) of the phenomena occurring on the Sun gives extraordinary opportunities to shed light on fundamental astrophysical questions, which are difficult to access from the study of more distant astrophysical objects. Thanks to its vicinity, the Sun can be studied in a detail by far impossible to reach in other stars. Besides, the Sun was the first star for which the traditional paradigm of merely positional astronomy was surpassed: prolonged and continuous time series are available for the obser- vations of some physical properties of the Sun (the systematic observation of solar spots dates back to four centuries ago, in the Western civilization, Bhatnagar and Livingston 2005). This role of solar astrophysics is particu- larly relevant for the phenomenon of magnetic activity, which represents the topic of the present work.

Since the ancient times, it was observed that the Sun presents inhomo- geneity and time variability, the solar spots being the most accessible and evident manifestation of the departure from a steady and homogeneous radiation. With the enormous increase and differentiation of observational ca- pabilities occurred in the last two centuries, the enormous broadening of the wavelength range observed, both from ground-based instruments and space- borne observatories, a large number of different structures and features in the atmosphere of the Sun was discovered, together with energetic phenomena and a temperature profile and consequent radiation emission which cannot be traced back to the conditions of radiative equilibrium. These processes are what is commonly referred to as “activity”.

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8 Chapter 2. The magnetic activity of late-type Stars

2.1.1 Large-scale structure of the solar atmosphere

Figure 2.1: Artistic view of the inner structure of the Sun and solar atmosphere’s layers. Courtesy NASA:

http://www.nasa.gov/images/content/462977main_sun_layers_full.jpg To understand the magnetic activity and its phenomenology, a basic knowledge of the solar atmosphere is needed. The atmosphere of the Sun (see Fig. 2.1) is defined as the volume surrounding the Sun whose lower layer is the region from which the Sun’s optical radiation was last scattered (optical depth τ = 1). This layer is called the photosphere, and is located at a radius of ≈ 10⁸m from the centre of the Sun. The outer atmosphere of the Sun extends above the photosphere, and is divided into two regions, roughly stratified in height: chromosphere and corona, separated by a so-called tran- sition region. At optical wavelengths, the photosphere outshines the fainter chromosphere and corona, and is visible as the solar disk emission, characterized by a black body spectrum with an effective temperature T_eff ≃ 5800 K.

The chromosphere is a region that extends above the photosphere for a few thousands km. In the chromosphere, the temperature decreases to a value of few thousands K from the photosphere up to ∼ 500 km, then rises slowly to ∼ 8000 K at 2000 km, and then rises sharply in its uppermost part to

∼ 25000 K (see Fig. 2.2). The chromosphere was discovered as a pink ring of emission at the solar limb that could be observed during solar eclipses (Lockyer, 1868). Its major feature is that cooling occurs mainly by radia-

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2.1. The Sun and the discovery of magnetic activity 9

Figure 2.2: Temperature as a function of height in the solar atmosphere for six diﬀerent models corresponding to six structures of the solar chromosphere observed at wavelength λ = 90 nm (A: a dark point within a convective cell;

B: the average cell center; C: the average quiet Sun; D: the average chromospheric network; E: a bright network element; F: a very bright network element) according to Vernazza et al. (1981).

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tion in strong resonant lines of abundant species such as Mg II and Ca II, Hα rather than in continuum, as is mostly the case in photosphere. Above the chromosphere, a shallow (∼ 100 km thick) transition region extends and connects the chromosphere with the corona. The transition region is characterized by a rapid increase of temperature from the values of the upper chromosphere to ∼ 1 MK. The corona is the outermost part of the solar atmosphere. As for the chromosphere, the corona was first observed during eclipses, as a bright extended non-homogeneous filamentary halo surrounding the obscured solar disk. Later, it was observed in optical coronal lines of highly-ionized species (Grotrian 1939, Edlén and Swings 1942). The corona was observed in radio wavelengths by Hey (1946), and then observed in X- rays during the first rocket flights (Burnight, 1949). The corona extends up to some stellar radii; on average, it is characterized by very low density (average electron density n_e∼ 10⁶−10⁹, up to 10¹¹cm⁻³), and high temperature, of some ∼ 10⁶K, that drops slowly at a distance larger than some stellar radii, where the corona fades into the stellar wind. However, the corona is a highly structured and non-homogeneous environment, as will be described in the following.

In the early studies, the solar atmosphere was regarded as a relatively static environment characterized by well-defined, homogeneous, gravitationally- stratified layers which lay one over the other without mixing. As observations progressively unveiled more and more details, a much more dynamical picture of the atmosphere become established, characterized by strong inhomogeneities and local structures, and dominated by highly dynamical processes and structures, in which a well-defined, static layer structure could no longer be recognised (see e.g. Aschwanden 2005), and this is particularly true for the corona.

2.1.2 Magnetic activity in the solar atmosphere

In contrast with a steady and homogeneous image of the solar atmosphere, we observe that magnetic fields cause inhomogenities, in the sense of uneven distribution of heat and plasma density, they determine impulsive release of energy, and heats up the whole atmosphere through various mechanisms. As mentioned, all these phenomena are what is generally referred to as magnetic activity. Magnetic fields can reduce the heat transfer and “devoid” regions from hot plasma (e.g., in the sunspots, through magnetic inhibition of convective heat transport) and, on the other hand, they may produce closed structures in which hot plasma is trapped. Besides, the energy stored in the magnetic field lines can be released to the surrounding plasma, causing impulsive heating. According to the particular region of the solar atmosphere, different manifestations of magnetic activity (small- or medium-scale structures shaped by the presence of magnetic fields and non-radiative-equilibrium

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Figure 2.3: Magnetic ﬁeld conﬁguration in correspondence of sunspots.

Sunspots are the photospheric footpoints of buoyiant magnetic ﬂux tubes.

The strong magetic ﬁeld inhibits the convective heat transfer from the stellar interior.

emission processes) are seen.

Solar photospheric activity

Solar spots are a typical example of how the magnetic field influences heat transport. They have been observed on the Sun since ancient times. In white light, they appear as darker spots on the solar photosphere, with a typical extension of ∼ 30000 km. They generally have a circular shape, but often irregular patterns are seen. Spots are interpreted as the footpoints of magnetic flux tubes buoyant from the deep solar interior and emerging from the photosphere (Fig. 2.3). Their lower brightness is due to the temperature, which is ≈ 1000−2000 K lower than the surrounding photosphere (≈ 5800 K).

This diﬀerence is allegedly due to the magnetic inhibition of convective heat transport in the deep regions of the spot. This model is strongly supported by the generally bipolar nature of sunspot groups (spots appear in couples of inverse magnetic polarity). The spot coverage (ﬁlling factor) on the Sun is very small, no higher than ∼ 0.5% of the total surface is covered by spots. The spot coverage is not constant, but it follows a cycle with a period of 11 years. During this cycle the spot coverage, and also all the other manifestations of magnetic activity (chromospheric and coronal) shift from

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Figure 2.4: Solar chromosphere with the most relevant small-scale structures (credits: A. Barber).

a maximum to minimum level and viceversa (Hale’s cycle, or “solar activity cycle”, Hale and Nicholson 1938).

Solar chromospheric activity

The images of the chromosphere taken at the wavelength of the predominant line emission show a complex web-like pattern, known as “chromospheric network”, characterized by alternating more and less bright regions and by different kinds of structures (Fig. 2.4), which outlines the spots seen on the solar photosphere, and evidences the presence of bundles of magnetic field lines that are concentrated by the fluid motion mainly on the border of spots. Near and around these regions, bright chromospheric emission regions, named “plages” (or “faculae”, but this term is generally referred to analogous photospheric bright regions) are observed, particularly bright in Ca and Hα lines. They appear to be associated with the concentrations of magnetic field occurring at spots. Regions of high magnetic field also harbour dark features, best observed in the Hα line, called “filaments”, or “prominences”:

dense and cooler clouds of material conﬁned within magnetic loops. Other structures, the so-called “spiculae”, are seen throughout the chromospheric network, as small jet-like eruptions appearing as short and dark streaks.

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Figure 2.5: Magnetic ﬁeld loops ﬁlled with hot plasma in the solar corona, observed by the TRACE observatory (Handy et al., 1999) at UV wavelengths.

Solar coronal activity

The solar corona presents an extremely rich zoo of diﬀerent structures and phenomena. First, it is necessary to distinguish three zones in the corona, whose relative size varies with time, corresponding to diﬀerent regions and structures in the underlying chromosphere and photosphere: the active re- gions, the quiet Sun and the coronal holes.

Active regions are the zones in which the most energetic phenomena are taking place. Their area is generally only a small fraction of the total surface, and they are located in regions of strong magnetic field concentrations, characterised by sunspot groups on the photosphere. Active regions are mainly made of closed magnetic field lines and subject to a wide number of different dynamic phenomena. They have the appearence of numerous loops filled with the plasma which remains confined by magnetic field lines to greater densities with respect to the medium surrounding the loops. Plasma-filled loops emit in soft X-rays and EUV (Fig.2.5).

The areas outside of active regions are historically referred to as the quiet Sun. Recent discoveries showed that also these regions harbour many dynamic processes. For these reasons, presently quiet Sun regions are de-

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Figure 2.6: The solar corona as seen during an eclipse. Coronal holes are clearly visible as regions in which the hot bright plasma becomes very thin in correspondence of the poles. Credits: NASA

ﬁned as all the regions of closed magnetic ﬁeld excluding active regions (see e.g. Aschwanden 2005).

The northern and southern polar zones of the Sun are dominated by open magnetic field lines (Fig. 2.6), that act as efficient conduits for flush- ing heated plasma coming from chromospheric upflows from the corona into the solar wind. Because of this efficient transport mechanism these regions are devoid of plasma most of the time, and thus they appear much darker than the quiet Sun regions, where heated plasma upflowing from the chromosphere remains trapped until it cools down and precipitates back. As a consequence, they appear as “holes” in the bright corona, and are called coronal holes.

Restructuring of the solar magnetic ﬁeld is associated with energetic events, called flares (see Sect. 2.1.2), which are impulsive release of mag- netic energy into the corona that causes rapid heating of the surrounding plasma and also involve the underlying chromosphere, yielding a rapid luminosity outburst in a wide range of wavelengths; The phenomenology of ﬂares and the underlying physical mechanisms require a more detailed description.

Flares

Flares are the results of impulsive energy release due to magnetic reconnec- tion between the coronal magnetic ﬁeld lines. The stellar dynamo (see Sect.

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Figure 2.7: 2D model of a magnetic reconnection process, driven by two oppositely directed inﬂows (in x–direction), which collide and create oppositely directed outﬂows (in y–direction) (Schindler and Hornig, 2000)

2.2) constantly generates new magnetic ﬂux from the tachocline ¹ which rises by buoyancy and emerges through the photosphere into the corona.

Convective motions and differential rotation in the interior of the Sun con- tinuously wrap up the magnetic field lines. As a consequence, the coronal magnetic field is constantly stressed and has to adjust by restructuring itself through topological changes. This occurs through the process of magnetic reconnection. In a simple 2-D model, when two oppositely directed magnetic field lines are pushed together, magnetic energy can build up between them.

The Lorentz force creates an electric field E₀ directed perpendicularly to the 2D-plane of the inflow magnetic field, and so a current sheet in the neutral layer, according to Ohm’s law. The current density increases as the sheet becomes thinner, until the initial configuration breaks, and one line connects to the other giving birth to two different magnetic field lines which are pushed outward (see Fig. 2.7), evolving to a lower-energy configuration, with the release of non-potential magnetic energy stored in the magnetic field lines (for a detailed description of the solar flare mechanisms see e.g. Aschwanden 2005).

Magnetic reconnection may be driven by diﬀerent geometric conditions.

The observation of flaring activity in the solar corona has led to the development of many models for flares that account successfully for the basic properties observed (see e.g. Aschwanden 2005). Particular geometries are possible on stars which possess an accretion disk or a surrounding envelope, as may occur in young stellar objects (YSOs). In Fig. 2.8 different magnetic

1The tachocline is the spherical surface which separates the inner radiative region, which rotates as a rigid body, from the outer, differentially-rotating radiative envelope in of solar-like stars. Physically, it is a thin layer characerized by a high rotational gradient.

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reconnection geometries that may occur on the solar and stellar surface or in YSOs are reported.

Magnetic reconnection has two immediate effects: it heats the plasma in the reconnection site, and it drives electron and ion acceleration. The first heating mechanism is the resistive or Joule plasma heating, in which the energy released by reconnection is dissipated by the Joule heating associated to the macroscopic currents. To be effective, due to the cooling time, this model requires anomalous resistivity. To coronal plasma heating may also contribute shocks occurring at the reconnection site. Also non-thermal accelerated ions can contribute to primary plasma heating, exciting strong kinetic Alfvèn waves (a class of magnetohydrodynamic waves, characterized by an oscillation of ions and electrons under the effect of an effective ten- sion aplied to the magetic field lines, Voitenko 1995, Voitenko et al. 2003).

Primary heating radiative signature comprises optical and UV continuum radiation; particles accelerated during the impulsive phase cause non-thermal radio emission, hard X-rays and γ-rays, (but this high-energy emission is generally outshined by the following chromospheric emission).

Nonthermal accelerated particles and thermal conduction fronts produced during the impulsive phase of the ﬂare, propagate downwards following the magnetic ﬁeld lines eventually hitting the transition region and the chromosphere, yielding secondary heating. In the chromosphere, electrons produce Bremsstrahlung emission with a direct signature in hard X-rays.

Chromospheric emission also comprises UV and optical wavelengths, in particular in the Hα and other Balmer lines, together with radio and microwave emission. The chromospheric heating determines the evaporation, or “ab- lation”, of chromospheric plasma at a temperature up to some 10 MK (in the Sun), that ﬁlls the magnetic ﬂux tubes, which becomes bright in the EUV and soft X-rays. Also radio bursts can be excited during this phase.

When thermal conduction and radiative losses prevail on the heating rate, the plasma starts cooling, emitting radiation at progressively longer wavelengths (Aschwanden, 2005).

Typical values for solar ﬂare luminosities span the range 26 < log L_X[erg/s]<

27 (0.5 − 8keV), with ﬂares up to log L_X[erg/s] ≃ 28. Flares as energetic as

∼ 10³²− 10³³erg have been observed (Aschwanden, 2005; Carrington, 1859).

The energetic phenomena and the structures described in this section are not a minor feature in the otherwise homogeneous and steady solar atmosphere. Instead, they represent its fundamental nature, they determine the structure itself of the various atmospheric environments (or layers), and de- fine them with their diversity. This is true both for photospheric structures (solar spots), and for chromospheric and coronal bright structures such as spiculae, plages, prominences, loop arcades, which are determined by the confinement of plasma in magnetic structures. The emission associated with bright magnetically-driven chromospheric formations (plages, spiculae) and corona (high-energy quiescent emission, flares) is due to excess heating al-

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Figure 2.8: Example of the configurations that may drive magnetic reconnection in the stellar coronae and in star-disk systems. The most simple situation (a), and the most popular solar flare scenario, is that in which a magnetic flux tube rises and expands. Its lower part shrinks, and the two sides come in contact, driving reconnection (e.g. Reeves et al. 2008). Several other possible geometric configurations are depicted in panels (b) through (f) (image from Benz and Güdel 2010).

legedly produced by magnetic processes. Diﬀerent mechanisms have been proposed to explain this excess of heat transfer to the chromosphere and

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corona, responsible for the inverted temperature gradient and for the very high temperatures observed, and consequenctly for the observed excess radiation, both magnetic, such as micro- and nano-flares, and mechanical, such as wave propagation, or combined mechanisms in which the magnetic field act passively allowing the propagation of waves, such as Alfven waves (Hall, 2008). What is certain, however, is that the magnetic field plays a fundamental role in the dynamic and radiative properties of the chromospheric and coronal structures.

2.2 The origin of solar magnetic field: the solar dy- namo

According to their scale and geometry, different kinds of magnetic fields can be distinguished on the Sun: global fields, in which the variation of the field parameters occur on a scale that is comparable with the solar radius; local fields, when the characteristic scale is much smaller than the solar radius, as in strong (∼ 1000 G) fields corresponding to sunspots. The global magnetic field of the Sun has two components, a poloidal and a toroidal one.

It is generally accepted that the magnetic ﬁeld of the Sun is sustained by an internal dynamo mechanism (Parker, 1955). The fundamental ingredients for the dynamo are the solar (diﬀerential) rotation and the convection.

Figure 2.9: Solar magnetic ﬁeld generation according to the αΩ dynamo mechanism (from http://www.konkoly.hu/solstart/stellar_activity.html).

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2.3. Magnetic activity on main-sequence stars 19

The fundamental hypothesis on which the dynamo theory relies is that in the solar interiors Ohm’s and Ampers’s law hold. Under this condition, it follows that a flow of electrically charged plasma can produce a magnetic field. In a simple, heuristic way, the standard dynamo mechanism (the so- called αΩ dynamo) can be described as follows (see Fig. 2.9). The rotating fluid of the convective envelope carries along the magnetic field lines of an initially dipolar configuration; under the effect of differential rotation, the field lines are dragged and amplified to form a stronger toroidal magnetic field. This is the Ω effect, that produces a toroidal field from a poloidal one.

The convective circulation produces an electromotive force that determines the onset of magnetic field loops in an axial plane. The resistivity of the fluid can make them reconnect and merge with other loops, rebuilding and amplifying a poloidal field (α effect). A detailed description of the dynamo mechanism can be found, e.g., in Dobler (2005); Charbonneau (2014).

2.3 Magnetic activity on main-sequence stars

2.3.1 Solar-like stars

As our Sun, the atmospheres of late-type main-sequence stars harbour a wide range of phenomena ascribed to magnetic activity, which pervade the phenomenology of their photosphere, chromosphere and corona. The detection and characterization of activity, and its connection with stellar rotation, give precious information on the physical properties, geometry and origin of stellar magnetic ﬁelds, and on the nature and distribution of internal dynamos across the Hertzsprung-Russell diagram.

When I talk about “solar-like” stars I refer typically to main-sequence stars of spectral type comprised between early-F and the limit of the fully convective regime (M3/M4-type), which are characterized by an outer convective layer and a radiative interior. However, the borders of these regions are not univocally deﬁned, as will be discussed in Sect. 2.3.3.

Starspots coverage As on the Sun, one of the most evident manifestations of activity is represented by spots. Photospheric spots are ubiquitous on main-sequence late-type stars. The spot activity observed on stars is often much more intense than on our Sun. Starspots can be directly observed with several techniques. Photometry, based upon the brightness variation due to the stellar rotation in the optical light curves, is the most common. Other techniques have also been used to envision starspots, including polarimetry, interferomety and micro-lensing produced during planetary transits.

Direct measurements of spot magnetic ﬁelds have been provided for G, K and M-type main sequence stars through line-broadening by, e.g., Valenti

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Figure 2.10: Example of spot cycles in the solar irradiance, and in V magnitudes for three active stars (Berdyugina, 2005).

et al. (1995), and through Zeeman-split spectral lines by Johns-Krull and Valenti (1996), Donati et al. (1997), Wade (2003). Among late-type stars, aside with stars showing a solar-like spot activity, or no observable spots, examples are found of stars presenting extremely large spot ﬁlling factors, up to ∼ 50% and higher (Strassmeier 1999, O’Neal et al. 2004). The prolonged observation of spot activity has also led to the detection of activity cycles on a time-scale of ∼ 10 yr, similar to that observed on our Sun (Fig. 2.10).

Chromospheric activity As described in Sect. 2.1.2, chromospheric activity is strongly correlated with the presence of starspots, since the regions characterized by the highest chromospheric activity (plages) tend to gather around spots, and can also be used as a proxy for the spot coverage. Chro- mospheric enhanced emission due to excess heating can be observed, on stars, in some spectral lines at infrared and optical wavelengths and some lines and continuum emission in the ultraviolet. The most relevant optical/infrared spectroscopic features that are used to this purpose are: the Ca II H&K lines (396.9 nm and 393.4 nm), the Ca infrared triplet (IRT) lines (848.9 nm, 854.2 nm, 866.2 nm), the Hα (656, 281 nm) and other Balmer lines, the Na I D₁, D₂ (589.6 nm, 588.6 nm), and Mg I b triplet (518.4 nm, 517.3 nm, 516.7 nm), the He I D₃ line (587.6 nm).

Large surveys of chromospheric spectral lines have been performed (e.g.

Wilson 1978,Duncan et al. 1991, Baliunas et al. 1995,Hall et al. 2007). Cro- mospheres appear to be ubiquitous among late-type stars. In general, the chromopsheric activity degree of solar-type stars, expressed as the ratio between H & K line brightness and bolometric luminosity (R_H&K), can be much

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higher than for the Sun (Baliunas et al., 1998). Solar-like cycles have been observed in chromospheric line intensity: Baliunas et al. (1998) found that a 60% of stars show periodic ﬂuctuations in spectral lines intensity which can be associated with a solar-like activity cycle, 25% show aperiodic modulations and 15% show a steady radiation.

Coronal activity and X-ray emission The best diagnostics for stellar coronal activity is represented by the high-energy radiation, in particular X-ray emission, which is basically associated with magnetic loops ﬁlled with hot emitting plasma. High-energy emission can be observed as steady, quiescent thermal emission or as impulsive outbursts in ﬂare events. Space-borne observatories (such as EINSTEIN, EXOSAT, ROSAT, XMM-Newton and Chandra) have provided large amounts of data on stellar X-ray emission.

Coronal emission is intrinsically multiwavelength, with another important diagnostic region located in the decimeter-centimeter radio wavelengths.

From various studies (Alcala et al. 1997, Schmitt and Liefke 2004), it emerges that the X-ray active coronae are ubiquitous along the whole low- mass main sequence. Since the ﬁrst surveys, late-type stars presented a wide spread in their coronal activity level: a spread of 2 − 3 orders of magnitude was observed by Pallavicini et al. 1981 in the X-ray/bolometric luminosity ratio of a sample of type G to M main-sequence stars, with overall X-ray luminosities up to at least 3 orders of magnitudes higher than the Sun; plasma with very high temperature, > 10 MK, is found especially in highly active stars (e. g. Swank et al. 1981). Generally, more active stars show higher coronal temperatures (e.g. Maggio et al. 1994, Schmitt 1997). Diﬀerent possible explanations have been given: for example, the increased activity determines directly hotter and denser features in the corona Maggio et al.

(1994), Güdel et al. (1997); alternatively, the increased activity leads to more numerous interactions between adjacent magnetic structures, thus increasing heating eﬃciency (Güdel et al. 1997, Audard et al. 2000.

All these evidence support the magnetic origin of the coronal radiation.

The direct correlation between the coronal X-ray emission and the magnetic ﬂux (Φ) was explored, e.g., by Pevtsov et al. (2003). They analysed spatially- resolved regions on the Sun and a sample of dwarf stars of spectral type G to M, with X-ray luminosity, and magnetic ﬂux. Fig. 2.11 shows the correlation between L_X and Φ for the solar structures and the sample of stars. The combined data reveal a single power-law relation (L_X ∝ Φ^1.13±0.05). The fact that the stellar coronae follow the same trend as the resolved solar surface suggests that the stellar coronae are heated by structures such as active regions that may be similar to those on the Sun.

Flares Flaring is the most evident manifestation of coronal magnetic activity. Solar-like ﬂaring activity appears to be ubiquitous among low-mass stars

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Figure 2.11: X-ray luminosity vs total unsigned magnetic ﬂux for solar structures and stellar objects. Dots: Quiet Sun. Squares: X-ray bright points.

Diamonds: Solar active regions. Pluses: Solar disk averages. Crosses: G, K and M dwarfs. Circles: T Tauri stars. Solid line: Power-law approximation of combined data set L_X ∝ Φ^1.15.

(e.g. Pallavicini et al. 1990). However, stellar flaring activity is characterised by a parameter space that is enormously larger than that of solar flares. The most productive spectral range for studying flares is the soft X-ray domain where the hot coronal plasma heated by the flare mechanisms radiates. Weak flares have been observed with soft X-ray luminosity of L_X≈ 10²⁶erg/s and soft X-ray total radiated energy E_X ≈ 1.5 · 10²⁸erg (Güdel et al., 2002), as well as extreme flares with peak luminosity L_X ≈ 10³²− 10³³erg/s and E_X ≈ 10³⁷ (e.g. Kuerster and Schmitt 1996, Getman et al. 2008). Electron temperature greater than 100 MK are common in stellar flares. Examples of extreme behaviour are seen also in the time profile, with rise phase lasting up to one day or longer, and decay time of order of one to several days (Tsuboi et al., 2016). The extreme flaring behaviour often observed is attributed to stronger magnetic fields, much higher filling factors, which result in larger total magnetic flux, and larger interaction volumes for the magnetic fields.

However, the limited instrumental sensitivity has restricted statistical studies essentially to highly active stars, in which a high cadence of luminous ﬂares was observed, so a bias may be present.

Optical and UV ﬂares have been widely observed on stars, as photometric brightness enhancement or spectroscopically resolved lines (mainly Ca, He and H Balmer lines, e.g. Hawley and Pettersen 1991,Hawley 2001). However, only a small fraction (∼ 10⁻⁴L_bol) of the ﬂare luminosity is due to optical

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emission (∼ 0.1 − 1 of X-ray energy, e.g. Kahler et al. 1982). A large search for optical ﬂares in a sample of G- to M-type stars was performed in the light curves provided by the Kepler satellite by Shibayama et al. (2013), which detected a large number of extremely intense “superﬂares”, and Davenport (2016).

Correlated X-ray and optical/UV flares have been observed. together with both X-ray flares without optical counterpart (e.g. Haisch et al. 1981) and viceversa (Doyle et al. 1986, Doyle and Byrne 1988). Due to the difficulty in coordinating multi-wavelength observations and to the intrinsically casual nature of flares, few examples of simultaneous optical/UV/X-ray flares have been reported.

2.3.2 Rotation/activity relation in solar-type stars

As described in Section 2.2, the main ingredients of the dynamo model are rotation and convection. Since the magnetic activity of solar-like stars is attributed to the presence of solar-like internal dynamos, the study of the correlation between the activity indicators and the stellar rotation is fundamental in order to verify the validity of this model, and eventually to characterize the dynamo model itself.

A correlation between magnetic activity and stellar rotation was found both for the photospheric, chromospheric and coronal activity manifestations. In order to provide the reader an insight into the issues of stellar rotation, which will also be useful to introduce the analysis performed in Chapter 4, I ﬁrst present here a brief outline of the most common diagnostics.

Rotation diagnostics

Several techniques are available to measure the rotation of stars: spectroscopic, photometric, interferometric and asteroseismologic. The most used spectrsoscopic techique is based on the observation of the broadening of photospheric line proﬁles due to the rotation of the stellar photosphere. If a star has a linear equatorial velocity v_e, the spectral broadening is given by

∆λL = (λ/c) · ve· sin i, where i is the inclination angle between the star’s rotation axis and the line of sight. The projected rotation velocity can be derived from the position of zeroes in the Fourier transform of the observed proﬁles. The v sin i technique is suited for fast rotators (v sin i ≥ 30 km/s), while it is not for slow rotators (v sin i ≤ 20 km/s), for which the noise is prominent.

The modulation in the overall optical brightness due to the uneven distribution of starspots on a not-resolved stellar rotating photosphere is one of the most used and reliable methods to measure the rotation rate. If we assume to have a single spot, or spot group on the photosphere of a star, we

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Figure 2.12: An optical light curve of a star (obtained by the Kepler pho- tometer, Sect. 3.2.1) showing a clear single-peak rotational modulation due to starspots. On top, a sketched view of a rotating star with a single spot group is plotted. The upward arrows represent the rotation axis. The sketch shows how the position of the spot group with respect to the line of sight is related to the observed ﬂux in the simple scenario of a single spot group.

observe, in the optical light curve, a ﬂux minimum at the phase at which the dominating (dark) spots are facing the observer, a maximum at the phase when the spots are on the hidden face of the star (Fig. 2.12). The main advantage of this technique is that it is free from geometric eﬀects, provid- ing the “true” rotation period. However, extensive photometric monitoring are needed in order to recover the periodic signal in the light curves (see e.g. Irwin and Bouvier 2009). Spots are not stationary; they appear and disappear, causing variations in the amplitude of the brightness modulation.

The photosphere may harbour more spot groups. This situation may yield a double-peak modulation in the light curve, which complicates the measure of the true period.

In case of differential rotation, spots located at different latitude evidence a different rotation period. Since the latitude of spots is expected to change during the activity cycle (as on the Sun), observations taken at different times will show different periods. Besides, if two spots groups at different latitude are present simultaneously, they will produce modulations with slightly different periods, and may cause beating in the light curve.

This eﬀect can be used to measure diﬀerential rotation, which is, however,

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Figure 2.13: Scatter diagram of X-ray luminosities versus projected rotational velocities for stars of spectral type G0-M5 (empty symbols) and F7- F8 (filled symbols), taken from Pallavicini et al. (1981). Different symbols represent different luminosity classes.

diﬃcult to detect and quantify.

Activity/rotation relation

Starspots coverage correlates with the rotation rate. The observation of overall optical brightness modulation showed, in particular, a general trend for which the spots modulation amplitude observed in optical light curves is inversely proportional to age and rotation period (Amado et al. 2001, Stelzer et al. 2003, Berdyugina and Järvinen 2005).

In a seminal work, Pallavicini et al. (1981) studied the correlation between coronal X-ray integrated luminosity, bolometric luminosity and projected rotational velocity (Fig. 2.13) for a sample of stars of various spectral classes using the data provided by the EINSTEIN observatory. This work came up with two fundamental ﬁndings: 1) a clear correlation was observed between the X-ray luminosity and the rotational velocity, which follows the relation L_X(erg s⁻¹) ∼ 10²⁷(v sin i [km s⁻¹])²), for the whole sample of stars of spectral type F7 to M5 including dwarf and giant stars. 2) Late-type

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stars do not show a correlation between X-ray luminosity and bolometric luminosity. It was also observed that F stars show only a weak correlation between X-ray luminosity and rotational velocity. Moreover, the ratio between the median X-ray luminosity and the median bolometric luminosity is signiﬁcantly higher than for the other spectral types, while the ratio between the median X-ray luminosity and the average rotational velocity is lower.

Another fundamental work was carried out by Pizzolato et al. (2003) in which the correlation between the X-ray luminosity/activity index and the rotation period/Rossby number (R₀²) was explored for a sample of ∼ 250 late-type dwarf stars. Noyes et al. (1984) introduced the Rossby number in the study of the activity/rotation relation. Dobson and Radick (1989) showed that this parameter correlates better with the X-ray activity than the rotation period, although this statement is debated (e.g. Reiners et al.

(2014)). The results (Fig. 2.14, 2.15) of this study present a clear bimodal trend in the relation between the X-ray luminosity and the rotation period for late-type stars, and the same holds for the X-ray activity index versus Rossby number relation; the X-ray luminosity/activity index presents a func- tional dependence (power law) on the rotation period/Rossby number up to a certain rotation rate (see Fig. 2.14, 2.15); for higher rotation rates, and corresponding lower Rossby numbers, saturation occurs, with X-ray luminosity and activity index that do no longer increase for increasing rotation rate (decreasing Rossby number). A similar bimodal trend was previously suggested by the results of Vilhu (1984) in both chromospheric and coronal activity. These results are consistent with the scenario of an internal dynamo which is responsible for the activity. Pizzolato et al. (2003) found a power law trend with index −2 in the correlation region between L_X (L_X/L_bol) and P_rot. For several mass (or B-V colour) classes they calculated the P_rot at which saturation sets in and the level of saturated L_X and L_X/L_bol, which depends on L_bol, consistently with the previous results.

The activity/rotation correlation, and in particular the bimodal nature observed in chromospheric and coronal activity is a strong clue that the observed activity originates from the magnetic fields produced by a rotation- powered dynamo : as the rotation rate increases, the dynamo gets more powerful, and the activity is higher; when the rotation rate reaches a certain threshold, the activity level saturates. Various effects have been ac- counted to explain saturation: limit in the flux that the dynamo can produce (Gilman 1983, Vilhu and Walter 1987); full surface coverage with active regions: consequently, the magnetic flux can no longer increase; disruption

2The Rossby number is a dimensionless quantity which is used to describe fluid flow in presence of roational motion. Here is defined as the ratio between the stellar rotation period and the convective turnover time, i.e. the timescale associated with the circulation inside a convective cell.

Ro=Prot

τ (2.1)

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Figure 2.14: X-ray luminosity-rotation relations from Pizzolato et al. (2003).

Left: X-ray luminosity vs. rotation period of field dwarfs (crosses) and cluster stars (squares). Leftward arrows indicate field stars with periods derived from v sin i data. Right: X-ray to bolometric luminosity ratio vs. rotation period for field dwarfs and cluster stars.

Figure 2.15: X-ray to bolometric luminosity ratio vs empirical Rossby number for all the stars in the sample of Pizzolato et al. (2003). The meaning of the symbols is the same as in Fig. 2.14.

of magnetic structures at fast rotation due to centrifugal forces (Jardine, 2004); centrifugally-induced eﬀective-gravity gradient from the equator to-

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wards the pole, which devoids equatorial regions from magnetic ﬂux (St¸epień et al., 2001).

2.3.3 Limiting cases of stellar dynamos

Among main-sequence late-type stars,the fully convective stars, of M3/M4 and later type, deserve particular attention. In a fully-convective stellar interior the classical αΩ dynamo cannot work, because of the lack of a tachocline.

Other dynamo mechanisms, such as the diffuse dynamo (α²Ω) can be relevant. According to this picture, a change in the magnetic flux across the stellar surface and in the topology of the magnetic field in the corona, determining a transition in the properties of the X-ray emission, may be expected around spectral class M3/M4. However, comparing the X-ray luminosity and activity index L_X/L_bol of stars below the fully-convective limit with the whole sample of M stars, a rather continuous trend emerges (e.g.: Johnson 1981,Drake et al. 1996, Fleming 2002). Flares have been often observed also in fully convective stars (Benz and Güdel 2010 and references therein).

At the opposite end of the main sequence, X-ray emission from early-type (O-B) stars is well known, but attributed to shocks produced in the stellar winds, and does not imply the presence of an active corona. A-type stars represent a particular case; they are not expected to produce strong enough winds to yield a significant X-ray emission from wind shocks; on the other hand, they have fully radiative interiors, so they are not expected to sustain a classical αΩ internal dynamo. Besides, even if A stars harbour magnetic fields, the lack of convective motion in the upper stellar layers would not al- low an effective transport of energy into the corona. However, X-ray bright emission was apparently observed also in early A-type stars Schröder et al.

(2008).

One hypothesis about this detection is that the X-ray emission was produced by unidentiﬁed later-type companion stars, in particular K and M stars. This scenario was particularly convincing for stars with high X-ray luminosity and hard spectrum (Golub et al., 1983). Besides, no clear rotation/activity correlation is observed for the X-ray emission from A-type stars, nor a correlation between rotation and the B-V colour (e.g. Panzera et al. 1999. However, more recent analysis of large samples of A-type stars has not led to an unambiguous attribution of the observed X-ray emission to companion stars (Schröder and Schmitt 2007, Stelzer et al. 2008). The evidence for ﬂares from A-type stars (Balona 2012, Balona 2015) are strongly debated. No consensus has been found yet, but it is possible that at least late-type A stars present phenomena to be ascribed to magnetic activity which indicates the existence of internal dynamos.

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2.4. Magnetic activity of Young Stellar Objects 29

2.4 Magnetic activity of Young Stellar Objects

The magnetic activity associated to the early phases in the life of late-type stars needs to be discussed separately: before positioning on the main sequence, indeed, pre-stellar objects pass through different phases, which need to be defined, and which may present significant differences in the activity level and in the underlying magnetic processes.

2.4.1 Evolution of Young Stellar Objects

While the formation of high-mass stars occurs in very short time, driving from a collapsing molecular nebula to the ignition of hydrogen in the newly formed object in ∼ 1−10Kyr, low-mass stars take some ∼ 1−10Myr of their life to evolve from the primordial cloud to the star positioned on the zero age main sequence, passing through different phases that are characterized by a different amount of circumstellar material with a different degree of interaction with the central forming star. These stages, from the collapse of a portion of a molecular cloud to the ignition of hydrogen in the star, are referred to as the Young Stellar Object (YSO) phase.

Different sub-classifications have been introduced according to the different evolutionary state of YSOs, referring to their physical and geometrical phenomenology or to their spectrum (Fig. 2.16). Lada (1987) proposed a classification based upon the shape of the Spectral Energy Distribution (SED) of the object, dividing YSOs in four classes (0, I, II, III), according to the slope of the SED in the infrared (α = dlog(λF_λ)/(dlogλ)), or equiv- alently to the bolometric temperature (T_bol) which reflects the presence or absence of envelopes and disks around the forming star. To these classes, another is often added, the “flat spectrum” class, which is intermediate between Class I and Class II. Another classification is based upon the strength of optical emission lines, and associates the evolutionary stage of the YSO with different classes of T Tauri stars (e.g., Feigelson and Montmerle 1999).

In the classiﬁcation by Lada (1987), Class 0 objects, corresponding to the earliest stage (duration: ∼ 10⁴yr), consist of a “seed” stellar nucleus, formed by the contraction of a molecular cloud, surrounded by an optically thick (A_V > few 100), cold (∼ 10 − 20 K, emitting at mm) gas and dust envelope with an extention of ∼ 10⁴AU, with the possible development of hot outﬂows. Class I objects are characterized by the persistance of an envelope, associated with a forming accretion disk of ∼ 10 − 100 AU in radius, through which the material in the envelope accretes onto the central forming star.

As a result of accretion, the surrounding envelope becomes optically thin, revealing the central object as a IR source. This phase has a typical duration of 10⁵yr. Later, the envelope disperses, settling into the disk. This phase corresponds to the Class II, or the classical T Tauri phase (Lada, 1987). At this stage, the star is fully convective, and contracts following the almost-

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Figure 2.16: YSOs classes: the spectral classiﬁcation by Lada (1987) is compared with the phenomenological classiﬁcation based on the evolution of the envelope and accretion disk (credits: André (2002)).

vertical Hayashi track (Hayashi, 1961) in the Hertzsprung-Russell diagram, staying at roughly the same surface temperature. Also the accretion disk

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subsequently vanish, leaving only the central objects, plus an eventual planetary system. At this stage (Class III), the star is referred to as a weak-line T Tauri star. At the end of the Hayashi track, the ignition of hydrogen fusion occurr; at this stage, stars with M_⋆ > 0.5M⊙ develope a radiative region, and move to the zero-age main sequence.

Objects in Class 0/I are generally referred to as “protostars”, and roughly correspond to the stages in the life of a YSO before the ignition of deu- terium fusion in the central forming star. Class II and Class III objects are instead referred to as “pre-main sequence stars”. Despite the usual discretiza- tion of these classes, there is a continuum of the circumstellar evolutionary phases, and thus a continuum in the phenomenological properties of YSOs.

The protostellar phase is diﬃcult to observe, because the star is still buried within its molecular cloud material. The information on what is going on in the central forming object can be achieved by the observation of the infrared and sub-millimeter emission, revealing information on the bolometric luminosity and spectrum of the embedded star and on their circumstellar material (Lada and Wilking 1984; Andre and Montmerle 1994; Andrews and Williams 2005).

2.4.2 Activity in pre-main sequence stars and protostars Many evidences for activity have been observed from YSOs, displying both similarity and differences with main-sequence solar-like stars. Pre-main sequence stars (PMSs), in particular, have been proven by extensive multiwavelength properties to be highly magnetically active. The characterization of the activity of protostars is instead more difficult and evidences more mixed. Among the activity phenomena in YSOs, coronal radiation, in the radio and X-ray band, chromospheric optical and UV emission lines, flaring activity and starspots are observed. Because of the intrinsic characteristics of YSOs not all the diagnostics for magnetic field and activity are usable. In particular, direct tracing of magnetic field and direct observation of starspots are difficult because of the high absorption; chromospheric activity indicators can be easily confused with accretion signature. Coronal X-ray emission (quiescent and flaring) is generally a good proxy for magnetic activity, even if deeply embedded objects may be subject to strong absorption of soft X-rays. The X-ray emission from YSOs consists of optically thin thermal bremsstrahlung and ionized metal emission lines from multitemperature plasmas with 1 < T_X < 100 MK (Montmerle, 1991).

Pre-main sequence stars

Classical and Weak-line T Tauri stars are allegedly considered to harbour internal dynamos. According to the standard dynamo model, to the high-

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est rotation rate of pre-main-sequence stars an enhanced level of activity (with respect to the main-sequence stars) corresponds Stelzer and Neuhäuser (2001), Preibisch et al. (2005).

Among weak-line T Tauri stars, several evidence of magnetic activity in the radio, optical and UV bands have been observed. Continuum radio emission was observed with speciﬁc luminosity three to six orders of magnitude brighter than powerful solar ﬂares (e.g. Garay et al. 1987, Chiang et al.

1996), together with radio flaring (Feigelson and Montmerle 1985, Phillips et al. 1991). In the optical and UV wavelengths, rotational starspot modulations, with filling factors up to 50%, have been observed (Rydgren and Vrba 1983, Strassmeier et al. 1994). Extremely energetic flares have been observed in Balmer lines (10³³− 10³⁴erg, Graham et al. 1995, Guenther and Emerson 1997). However, in non-X-ray wavebands, evidence of magnetic activity for Classical T Tauri stars are quite difficult to identify. Few detections in the centimeter have been found, while in the optical band the manifestations of activity are often covered by star-disk interactions.

The X-ray emission of Classical and Weak-line T Tauri stars does not present significant differences. Both classes of stars are widely detected in X-rays. Quiescent sof X-ray luminosities range between 10^28.5− 10³¹erg/s (Feigelson and Montmerle, 1999). X-ray luminosity scales with bolometric luminosity, with L_X/L_bol ≈ 10⁻⁴, which is below the saturation level seen among late-type evolved stars (10⁻³). The X-ray surface flux increase with age, as the star contracts; this probably means that the active regions oc- cupy a roughly constant area as the star contracts. L_Xis also correlated with age, effective temparture, mass and rotation rate (although the correlation is noisy), as for evolved late-type stars (e.g. Feigelson et al. 1993). X-ray flaring activity is ubiquitous in PMSs. Studies on the solar-type PMSs in the Orion Nebula Cluster have been performed by, e.g., Wolk et al. (2005) (COUP project). The flaring activity of YSOs presents many similarities with the flaring activity on the Sun and on sun-like magnetically active stars.

However, it is characterized by events that are orders of magnitudes more energetic than the flares observed on the Sun and have an high energy also compared to active Sun-like stars, both in the radio, optical and the X-ray band (Stelzer and Burwitz 2003, (Maehara et al., 2012)). The X-ray luminosity function is broad, spanning 28 < log L_X[erg/s]< 32 (0.5˘8keV), with a peak around log L_X [erg/s]∼ 29 (Feigelson et al., 2005), and logarithm of the integrated energy log E_X[erg]≃ 34−36, up to log E_X[erg]≃ 37 (Preibisch et al., 1995). Very high plasma peak temperatures have also been observed in flares on YSOs in ONC, with typical values up to 100 MK (Favata et al., 2005). These values are far above those observed in solar flares and hogh also compared to flares observed on solar-like stars. Flares from PMSs are not only brighter and more energetic, but they also show higher frequency than observed on the Sun and late-type main-sequence stars (Güdel 2004 and references therein). The COUP flares show a frequency ∼ 100 times higher

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than the most powerful ﬂares observed in the contemporary Sun, with about one ﬂare a week with L_X > 2 · 10³⁰erg/s.

Protostars

Since protostars are deeply embedded in thick envelopes (A_V ≈ 10 − 50 in Class I, up to A_V ∼ 1000 for Class 0), also detecting X-ray emission is intrinsecally diﬃcult, and the knowledge of their magnetic activity is frag- mentary. In mid-1990s, unequivocal X-ray detections of Class I stars emerged (Casanova et al. 1995,Imanishi et al. 2001b). The observed L_X are close to saturation (L_X/L_bol ≈ 10⁻³, Ozawa et al. 1999); extreme values have been reported (L_X ≈ 10^32.5− 10^33.5, e.g. Preibisch 1998). Very high characteristic temperatures (20 − 40 MK, Tsujimoto et al. 2002,Imanishi et al. 2001a) are observed. The observations of X-ray emission from Class I protostars remain however quite rare. The situation is more drastic for Class 0. It is not clear if Class 0 protostars harbour magnetic activity, and in particular X-ray emission, and if they have internal dynamos, and if so, how they work. The classiﬁcation is often critical. A report of hard X-ray emission from two Class 0 protostars in the OMC 2/3 region was presented by Tsuboi et al. (2001), but other researchers classify these objects as Class I (Niel- bock et al., 2003). Contrary to the few possible detections of X-ray emitting Class 0 protostars, a large number of Class 0 objects have been observed, e.g., with Chandra and other X-ray observatories, e.g. in the Serpens and NGC 1333 embedded clusters Getman et al. 2002, Preibisch 2004) ρ Ophiuchi Cloud complex (Preibisch, 2004), without being detected in X-rays.

Flares observed on Class I protostars may be extremely intense. Several X-ray ﬂares have been observed from Class I stars, many of which exceed- ingly large, with total soft X-ray energy up to ≈ 10³⁷erg (Koyama et al.

1996, Preibisch 2003). A comparison between X-ray flares from Class III and Class I stars by Imanishi et al. (2003), based upon MHD models, showed that Class I stars tend to have stronger flares, realistically requiring larger volumes, thus suggesting extended star-disk magnetic fields as the source of flares (e.g. Grosso et al. 1997).

All these evidences suggest that YSOs (at least pre-main-sequence stars and Class I protostars) present an enhanced solar-type magnetic activity and harbour intense surface magnetic ﬁelds. This is not surprising in rapidly rotating cool stars with deep convective zones, as the YSOs are. However, the standard αΩ solar dynamo model cannot work in fully-convective YSOs, so the comparison with solar activity is based more on phenomenological similarity than on physical arguments.

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Chapter 3

Database

In the present work I make wide use of data taken from the database of the XMM-Newton and Kepler missions. In the analysis of the XMM-Newton data, I make use of the pipelines for the characterization of X-ray time variability and of the associated products provided by the EXTraS Project.

In this chapter, a brief outline of the two missions, the EXTraS Project and their databases is given.

3.1 XMM-Newton data

The study of the X-ray emission from stars and Young Stellar Objects presented in this work is based upon the analysis of the X-ray observations performed by the XMM-Newton mission. XMM-Newton provides high per- formance X-ray photometry and spectroscopy together with wide and deep sky coverage.

X-ray Multi-Mirror (XMM)-Newton is a space-borne X-ray observatory launched by the European Space Agency (ESA) in 1999, and still fully oper- ative. The aim of XMM-Newton is to study the galactic and extragalactic sources of X-rays, in order to explore the physics of the high-energy emission processes, in particular to perform sensitive medium-resolution spectroscopy and broad band imaging spectroscopy. The concept at the basis of the de- sign of XMM-Newton was to create an X-ray observatory which combined diﬀerent multi-mirror X-ray telescopes (Fig. 3.1).

3.1.1 On-board instruments

XMM-Newton carries two diﬀerent types of telescopes: three multi-mirror grazing-incidence Wolter type-1 X-ray telescopes for imaging and spectroscopy (main instrument), and one additional optical monitor (OM) for the observation in the optical (visible and UV) range. Each of the X-ray telescopes consists of 58 grazing incidence composite parabolic-hyperbolic coaxial mir-

35

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36 Chapter 3. Database

Figure 3.1: Schematic of the instruments onboard XMM-Newton satellite.

rors with a common focus, which provide a high effective area over a wide photon energy range, from less than 1 keV to more than 10 keV. Each telescope has a diameter of 70 cm (largest mirror) for a focal length of 7.5 m, and is provided with a baffle for protection against visible and X-ray straylight and a deflector for soft electrons. At the primary focus of each telescope there is one of the three CCD devices that compose the European Photon Imaging Camera (EPIC), one pn-junction camera (Strüder et al., 2001) and two MOS (Metal-Oxide-Semiconductor) cameras (MOS1 and MOS2, Turner et al. 2001). EPIC allows to perform sensitive imaging observations over a wide field of view and energy range, with moderate spectral, angular and timing resolution. The field of view of each EPIC detector is ≈ 15 arcmin of radius (Fig. 3.2). The CCD of the EPIC cameras are energy sensitive, enabling non-dispersive spectroscopy. The EPIC camera can detect photons with high efficiency in the range 0.15 − 15 keV. The combined effective area of the three EPIC instruments is plotted in Fig. 3.5 The spectral resolution depends on the intrinsic energy resolution of the CCDs pixels, and is a function of energy.The angular resolution of the EPIC instruments depends on the pixel size in the CCD of the detector and on the point spread function (PSF) of the X-ray telescopes. The projected angular size of each pixel of the pn CCDs 4.1 arcsec; for the MOS CCDs it is 1.1 arcsec. The full width at half maximum (FWHM) of the PSF for the pn and MOS, at on-axis position (near the centre of the detector) is respectively < 12.5 and 4.4.

The EPIC detectors can be operated in diﬀerent modes, depending if the purpose of the observation is to perform the imaging of the largest possible ﬁeld of view or to achieve the highest possible timing resolution. Four ob- serving modes are available for the MOS cameras and six for the pn camera.

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3.1. XMM-Newton data 37

Figure 3.2: Field of view of one of the EPIC MOS cameras (MOS1, left, in “Full Frame” observation mode ) and of the EPIC pn camera (right, in

“Extended Full Frame” mode).

Figure 3.3: Combined eﬀective area of the EPIC pn/MOS1/MOS2 detectors.

The different lines represent the three different filters (thin, medium, thick) which can be used to protect the detector from IR/optical photons.

In the standard imaging mode (Full frame) the time resolution is 2.6 s for the MOS and 73.4 ms for the pn. The highest time resolution (non-imaging mode) for the MOS is 1.75 ms and for the pn is 0.03 ms in timing mode and 7 µs (but with the very low duty cycle of 3%).

The two telescopes read out by the MOS cameras are also equipped with